adsabs/FOCAL huggingface dataset:
The ultraviolet (UV) upturn or excess is an unexpected rise in flux in the spectral energy distributions (SEDs) of early-type galaxies shortwards of 2500 Å. It appears to be a nearly ubiquitous property of spheroids and bulge-dominated galaxies (see e.g. Yi 2008, 2010 for a recent review), although their generally old, metal-rich, and quiescent stellar populations (e.g. Thomas et al. 2005, 2010) should contain no sources capable of providing significant flux below the 4000 Å break. While many candidates have been proposed, the source population is generally agreed to consist of hot horizontal branch (HB) stars (Greggio & Renzini 1990; Brown et al. 1998a) and it now appears most likely that this population is helium-rich and formed in situ at high redshift (Ali et al. 2018a,b,c). While such stars are now known to exist in globular clusters in our Galaxy (e.g. see Norris 2004; Piotto et al. 2005, 2007) and likely elsewhere (Kaviraj et al. 2007b; Mieske et al. 2008; Peacock et al. 2017), the origin of such high helium abundances is unclear.1 It has been suggested that the effect may depend on the environment: stratification of helium in the centres of clusters might create populations of galaxies with high helium abundances (Peng & Nagai 2009). On the other hand, it is also possible that the helium enrichment depends on the details of early star formation, as in globular clusters. Differences in the originating mechanisms or time-scales may be reflected in different evolutionary histories for galaxies in clusters and the field. Atlee, Assef & Kochanek (2009) find a slow decrease in the strength of the upturn for a sample of bright field elliptical galaxies at 0  z  0.65 compared to the nearly constant colour for very bright ellipticals in the works of Brown et al. (1998b, 2000, 2003) and brightest cluster galaxies (Loubser & Sánchez-Blázquez 2011; Boissier et al. 2018). In our work on much larger samples of cluster early types with luminosities down to L*, we also see no evolution at z  0.55 (Ali et al. 2018a,b,c) but then detect a rapid reddening in the UV colour at z = 0.7; this may be consistent with the last data point in Atlee et al. (2009), despite the large errors and small number statistics. Le Cras et al. (2016) also find evidence for evolution in the UV upturn at z > 0.6, albeit from stacked spectra and using spectroscopic indices sensitive to the HB population, for a sample of very luminous BOSS galaxies (as opposed to the UV photometry used by other studies).
Citation: Ali et al. 2018a
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Citation: Ali et al. 2018a
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It may be appropriate to single out at this point the recent and interesting works by Lin & Ishak (2017a,b) in which the authors run a so-called (dis)cordance test based on using a proposed index of inconsistency (IOI) tailored at finding possible inconsistencies/tensions between two or more data sets in a systematic and efficient way. For instance, it is well known that there is a persistent discrepancy between the Planck CMB measurements of H0 and the local measurements based on distance ladder (Riess et al. 2016, 2018b). At the same time, if one compares what is inferred from Planck 2015 best-fitting values, the LSS/RSD measurements generally assign smaller power to the LSS data parametrized in terms of the weighted linear growth rate f(z)σ8(z). This feature is of course nothing but the σ8-tension we have been addressing in this paper. It is therefore natural to run the IOI test for the different kinds of H0 measurements and also to study the consistency between the H0 and the growth data. For example, upon comparing the constraints on H0 from different methods, Lin & Ishak (2017b) observe a decrease of the IOI when the local H0 measurement is removed. From this fact they conclude that the local measurement of H0 is an outlier compared to the others, what would favour a systematics-based explanation. This situation is compatible with the observed improvement in the statistical quality of the fitting analysis by Solà, Gómez-Valent & de Cruz Pérez (2017b,c) when the local H0 measurement is removed from the overall fit of the data using the RVM and the ΛCDM. In this respect, let us mention that a recent model-independent analysis of data on cosmic chronometers and an updated compilation of SNIa seem to favour the lower range of H0 (Gómez-Valent & Amendola 2018), what would be more along the line of the results found here, which favour a theoretical interpretation of the observed σ8 and H0 tensions in terms of vacuum dynamics and in general of DDE (cf. Fig. 10).
Citation: Riess et al. 2016; Solà, Gómez-Valent & de Cruz Pérez (2017b
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Citation: Lin & Ishak (2017b)
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Citation: Lin & Ishak (2017; b
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These results show that the shape of the mass function and the position of the mass peak of massive clusters have little evolution over the course of the galaxy collision of more than 1 Gyr. We note that the evolution of SCs is subjected to several destructive processes (e.g., Gnedin et al. 1999a; Fall & Zhang 2001). For low-mass clusters (





), the destruction is mainly dominated by two-body relaxation processes, in which the mass of a cluster linearly decreases with time until it is destroyed. For more massive clusters, the evolution is primarily influenced by stellar evolution at early times (





 Myr) and by gravitational shocks at later times. These effects are included in our hydrodynamic simulations, but the mass resolution is not high enough to resolve the processes realistically, since the two-body relaxation and stellar evolution depend on individual stars. However, the cluster disruption timescale due to two-body relaxation is proportional to the cluster mass, trlx ∼ 1.7 Gyr × (Mclus/104 





)0.62 × (T/104 Gyr−2)−0.5, where Mclus is the cluster mass and T is the tidal strength around the clusters (Kruijssen et al. 2012). Using the minimum cluster mass of 





 in our simulation and a typical range of tidal strength in the nuclear region (since most of these clusters are concentrated around galaxy nuclei) of 





 (Renaud 2010), we find that the range of the disruption timescale is 





 Gyr. For more massive SCs, the timescale is even longer, beyond our run time of 1 Gyr. Therefore, the two-body relaxation may not have a major disruptive effect on these SCs. Furthermore, the mass loss timescale due to gravitational shocks (tsh) depends strongly on the cluster density, 





 (Kruijssen et al. 2012). The majority of our clusters have a density range of 





, as shown in Figure 9, which suggests 





 Gyr, much longer than the Hubble time. So gravitational shocks may not have a significant impact on the clusters in our simulation. In addition, as demonstrated by Renaud & Gieles (2013), SCs formed in galaxy mergers are also affected by the intense tidal field of the galaxies, more so for clusters in the merger remnant than for the ejected ones. We note, however, the clusters in their simulations have masses 





, and mass loss decreases as the cluster mass increases. For example, for a cluster to increase its mass by a factor of 2, from 





 to 





, its survival rate (fraction of initial mass survived) after 1 Gyr increases from 0.6 to 0.7. Extrapolating this trend to our clusters, which are ∼10 times more massive than those in Renaud & Gieles (2013), we argue that destruction from tidal fields likely has negligible effects on the clusters we consider here.
Citation: Kruijssen et al. 2012
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Citation: Kruijssen et al. 2012
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The second significant source of energy dissipation is stream–stream collisions near apocenter from GR apsidal precssion. As in Dai, McKinney & Miller (2015) and Bonnerot et al. (2017), we calculate the ingoing collision stream velocities by finding the intersection point of two elliptical orbits, shifted by an instantaneous pericenter shift δϖ from GR apsidal precession. Since an eccentric orbit undergoes apsidal precession at a rate (assuming af ≫ Rg)
(24)$$\begin{eqnarray*}
\left. \frac{\partial \varpi }{\partial t} \right|_{\rm GR} = \frac{3 R_{\rm g}n}{a (1-e^2)},
\end{eqnarray*}$$we have after one orbital period
(25)$$\begin{eqnarray*}
\delta \varpi \simeq \frac{3 \pi R_{\rm g}}{r_{\rm p}} = 11.5^\circ \, \beta \frac{\bar{M}_\bullet ^{2/3} \bar{M}_\star ^{1/3}}{\bar{R}_\star }.
\end{eqnarray*}$$With δϖ, one can calculate the relative stream velocities during the collision, leading to a loss of energy of (Dai et al. 2015; Bonnerot et al. 2017)
(26)$$\begin{eqnarray*}
\delta E_{\rm ss} &amp;=&amp; \frac{e_{\rm f}^2 G M_\bullet }{2a_{\rm f}(1-e_{\rm f}^2)} \sin ^2 \left(\frac{\delta \varpi }{2} \right) \simeq \frac{G M_\bullet }{4 r_{\rm p}} \sin ^2 \left(\frac{\delta \varpi }{2} \right),
\end{eqnarray*}$$hence
(27)$$\begin{eqnarray*}
\mathcal {V}_{\rm ss} = \frac{\delta E_{\rm ss}}{E_{\rm f}} = \frac{a_{\rm f}}{2 r_{\rm p}} \sin ^2 \left(\frac{\delta \varpi }{2} \right).
\end{eqnarray*}$$When δϖ ≪ 1, equation (27) reduces to
(28)$$\begin{eqnarray*}
\mathcal {V}_{\rm ss} \simeq \frac{9 \pi ^2}{16} \frac{R_{\rm g}^2 a_{\rm f}}{r_{\rm p}^3} = 0.125 \, \beta ^3 \frac{\bar{M}_\bullet ^{5/3} \bar{M}_\star ^{1/3}}{\bar{R}_\star ^2} \left(\frac{t}{t_{\rm f0}} \right)^{2/3}.
\end{eqnarray*}$$Different sources of energy dissipation need to be compared with the total loss of energy required to completely circularize the disc. If the debris circularizes with constant orbital angular momentum, the final semimajor axis is given by $a_{\rm circ} = (1 - e_{\rm f}^2)a_{\rm f}\simeq 2 r_{\rm p}$. Hence the orbital energy of a completely circularized debris stream is given by
(29)$$\begin{eqnarray*}
E_{\rm circ} = - \frac{G M_\bullet }{4 r_{\rm p}},
\end{eqnarray*}$$so the fractional binding energy which must be lost to circularize the stream is
(30)$$\begin{eqnarray*}
\mathcal {V}_{\rm circ} = \frac{E_{\rm circ}}{E_{\rm f}} = \frac{R_{\rm t}^2}{4 r_{\rm p}R_\star } \left(\frac{t}{t_{\rm f0}} \right)^{2/3} = \frac{25}{\beta } \frac{\bar{M}_\bullet ^{1/3}}{\bar{M}_\star ^{1/3}} \left(\frac{t}{t_{\rm f0}} \right)^{2/3}.
\end{eqnarray*}$$We see that $\mathcal {V}_{\rm circ} \gg 1$ for typical parameters.
Citation: Dai et al. 2015
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In this Letter we propose to address these apparent difficulties and account for many of the observed differences between the properties of planets in cool and hot stars by postulating that, in addition to the tidal interaction with existing close-in planets, a large fraction of the stellar hosts—both cool and hot—ingest a hot Jupiter early on in their evolution. This proposal is motivated by the expectation that a large fraction (up to 80% according to Trilling et al. 2002) of solar-type stars possess giant planets during their pre-MS phase, and that a large fraction of the giant planets that form in a protoplanetary disk on scales 





 migrate close to their host star before the disk is dispersed (e.g., Ida & Lin 2004). Numerical simulations incorporating an N-body code and a 1D α-viscosity disk model (Thommes et al. 2008) demonstrated that this behavior can be expected for disks with 





 that are sufficiently massive. The inward planet migration is likely stopped by the strong (





 kG) protostellar magnetic field that truncates the disk at a radius (rin) of a few stellar radii (e.g., Lin et al. 1996). Gravitational interaction with the disk causes a planet reaching rin to penetrate into the magnetospheric cavity and, if it is massive enough, to undergo eccentricity excitation that can rapidly lead to a collision with the star (e.g., Rice et al. 2008). It was, however, inferred that if Mp is sufficiently small (





), the planet would remain stranded at a distance where its orbital period is ∼0.5 of that at rin until well after the gas disk disappears (on a timescale of ∼ 106–107 years). In our proposed scenario, the primordial disk orientations span a broad angular range that is reflected in the orbital orientations of the stranded planets. When the latter are ingested by tidal interaction with the host star (on a timescale 





 Gyr), the absorbed angular momentum is sufficient to align a solar-mass star in that general direction, but not an MS star with 





 K. This is because cool stars have significantly lower angular momenta at the time of ingestion than hot stars as a result of a more efficient magnetic braking process and of a lower moment of inertia. Given the proximity of the stranded HJs (SHJs) to their host stars (





days), they can be expected to have been ingested by the time their parent planetary systems are observed; however, giant planets farther out can continue to interact with their host stars and potentially affect their measured obliquities. In our simplified formulation, we model the SHJs using as parameters their characteristic mass MSHJ and the fraction p of systems that initially harbored an SHJ. By comparing the predictions of this model with the observational data, we infer 





 and 





.
Citation: Ida & Lin 2004
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Citation: Thommes et al. 2008
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Low-frequency observation of radio sky is also essential for the study of the astrophysics at play in various evolutionary stages of different Galactic and extragalactic sources. Radio emission at low frequencies, together with their redshift information, can be used to infer several astrophysical properties associated with the sources. In general, the source distribution is assumed Poissonian (with the possibility of clustering following single power law) (Ali et al. 2008; Jelić et al. 2010; Trott et al. 2016). Source counts are also modelled via a single power-law distribution (Hurley-Walker et al. 2016; Intema et al. 2017; Franzen et al. 2019). However, several recent studies have shown deviation from the single power-law model (Williams et al. 2016; Prandoni et al. 2018; Hale et al. 2019). Thus more detailed studies both in wide field as well as deep fields are required for generating a fiducial model of sources at low frequencies. The differential source counts at these are also useful for constraining the nature of sources. At frequencies in the GHz range and upwards, the source properties are well characterized. However, at lower frequencies, there is a lack of consensus for the same. It has been seen in previous studies that at these frequencies, AGNs dominate in the flux density scales down to few 100 μJy while SFGs and radio quite AGNs become dominant below ∼100 μJy (Simpson et al. 2006; Mignano et al. 2008; Seymour et al. 2008; Smolčić et al. 2008; Padovani et al. 2009, 2011, 2015; Prandoni et al. 2018; Hale et al. 2019). This is inferred from the flattening of the source counts below 1 mJy. However, such studies are very few on account of the limitations in reaching the required SNR. Thus empirical constraints at frequencies ≲1.4 GHz are limited. Therefore the study of the low-frequency radio sky is vital for fiducial modelling of foregrounds for 21-cm cosmology as well as constraining the physics and the astrophysics of the sources.
Citation: Prandoni et al. 2018
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Citation: Prandoni et al. 2018
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Citation: Jelić et al. 2010
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Applying the P&T profiles defined in Sect. 2 for the three regions of IRAS 16293 our computations imply that 56 mechanisms possess valid reaction thermodynamic function values to proceed with the formation of H2CO. Twenty-six of these can produce H2CO in region I, 32 in region II, and 33 in region III (please see Table 7). These are the amended reactions with detected reagents in every region of the disk. The presence of methane (CH4) and neutral atomic carbon (CI) in this astronomical object, in addition to trans-HONO in region III (detected at ~ 150 AU from the core – Coutens et al. 2019), confirms the hypothetical feasibility of these mechanisms. Moreover, because nitrosyl hydride (HNO) has not been detected in IRAS 16293, R84 cannot be included in any circumstellar envelope region, as it was in DCDMCs. In Fig. 3 we show the abundance of reactants and formaldehyde itself already detected in the three regions of the disk. They have been calculated for the methylene (⋅ 3CH2) (Vasyunin & Herbst 2013), methyl radical (⋅CH3) (Sakai et al. 2012; Woon 2002), and methane (CH4) (Sakai et al. 2012) ground electronic state abundances rather than detected. The reason is that they are planar (or linear) symmetric-top molecules and therefore have a small permanent electric dipole moment, which makes them undetectable through terrestrial radio telescopes by rotational spectroscopy. On the other side of the spectrum, neither high-resolution infrared absorption nor ultraviolet-visible spectroscopy telescopes can identify them in the gas-phase because they lack of a powerful enough source of electromagnetic radiation. The abundance values for the triplet ground-state of methylene (⋅CH2(X3B1)) have therefore been estimated by using Monte Carlo algorithms (Vasyunin & Herbst 2013) which in turn are based on known abundances of derived species. The presence of the methyl radical (⋅CH3) in region I just before the CO depletion zone from icy grains should be negligible because ⋅CH3 is thought tobe mainly produced in this region by CH4 photodissociation (Sakai et al. 2012) and CH3OH (Woon 2002) desorbs from the grains in region II. Consequently, reactions based on ⋅CH3 were considered only for regions II and III, assuming for both an abundance of 6.40 × 10−9 cm−3 (Sakai et al. 2012).
Citation: Sakai et al. 2012
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Citation: Sakai et al. 2012
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Citation: Sakai et al. 2012
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Citation: Sakai et al. 2012
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Fukui et al. (2014, 2015) compared the Planck dust data and the gas emission data such as 





 and 





 for the solar neighborhood. The 





 emission is generally assumed to be optically thin, as written in textbooks. If the dust properties are uniform and DGR is a constant, it is expected that the velocity-integrated intensity of the 





 spectrum (





) is proportional to 





 for the data points where the 





 emission is not detected. Fukui et al. (2014, 2015), however, found that the correlation between them is not so good. By introducing 





 into the 





 correlation plot, these authors discovered that the poor correlation in the 





 plot is mainly due to the data points where the density is high and 





 is low. There are two main possibilities to explain this bad correlation between 





 and 





 one is the presence of optically thick 





 gas, and the other is the presence of “





-dark 





 gas,” which is 





 gas without the 





 emission (e.g., Wolfire et al. 2010; Planck Collaboration et al. 2011c; Langer et al. 2014). Fukui et al. (2014, 2015) investigated the 





 fractions in the hydrogen gas by referring to the UV measurements (Gillmon et al. 2006), and found that the fractions are typically 





 toward the lines of sight whose column densities are up to at least 





.1

1
Note that in Figure 16 of Fukui et al. (2015), the 





 fractions are somewhat larger than 10% toward two Galactic B-type stars, HD 210121 and HD 102065 (Rachford et al. 2002). These two stars may be contaminated by their own localized gas, and therefore the 





 fractions for the local ISM are possibly not reliable toward them. For this reason, Fukui et al. (2017; see Section 4.3) did not use the results obtained in Rachford et al. (2002).
 That is, 





 dominates 





, and the “





-dark 





 gas” would not be a dominant component in the local ISM. Therefore, these authors concluded that there exists a large amount of optically thick 





 gas in the local ISM, whose typical optical depth is 





, and the amount of the 





 gas is underestimated by 





 if a correction for the opacity effect is not applied. In Fukui et al. (2014, 2015), these analyses were made for the high Galactic latitude at 





 where the gas density is low. An open issue discussed in Fukui et al. (2014, 2015) is if 





 obeys a simple linear relationship with the 





 or not. A study in the Orion A molecular cloud (Roy et al. 2013) indicates that the dust optical depth is proportional to the 1.28th power of 





, rather than a simple linear relation. If correct, this nonlinearity may be ascribed to the dust evolution at high column density. This suggests that it is possible to apply a minor modification of the method of Fukui et al. (2014, 2015) in order to improve the accuracy in 





.
Citation: Gillmon et al. 2006
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Half a century ago, a series of theorems laid the ground for the Kerr hypothesis (Israel 1967; Carter 1971; Robinson 1975); according to these no-hair theorems, the only stationary, axisymmetric, asymptotically flat, regular outside of the horizon solution to four-dimensional GR when the matter fields feature the same isometries as the spacetime is the Kerr BH. Notwithstanding their significance, there are many ways with which to circumvent them and discover different solutions. Still, in four dimensions, hairy BHs have been described in different theories of gravity, such as Einstein–Yang–Mills (Bizon 1990; Künzle & Masood-ul-Alam 1990; Volkov & Galtsov 1990; Breitenlohner et al. 1992; Kleihaus & Kunz 1998, 2001; Kleihaus et al. 2004), scalar-tensor (Bocharova et al. 1970; Bekenstein 1974; Kleihaus et al. 2015; Collodel et al. 2020a), and Gauss–Bonnet theories (Kanti et al. 1996; Kleihaus et al. 2011, 2016; Antoniou et al. 2018; Doneva & Yazadjiev 2018; Silva et al. 2018; Cunha et al. 2019; Collodel et al. 2020b; Herdeiro et al. 2021; Berti et al. 2021). Remarkably, by dropping the assumption that the matter fields must be stationary and axisymmetric, Herdeiro and Radu found solutions in the context of GR where BHs have hair (Herdeiro & Radu 2014b, 2015), by minimally coupling to gravity a complex scalar field that depends on time and on the axial coordinate while its energy-momentum tensor still possesses the respective isometries; see Herdeiro et al. (2015, 2016a, 2016b), Brihaye et al. (2016), and Delgado et al. (2016) for generalizations. These are known as scalarized Kerr black holes (KBHsSH) and they are the object of study of this paper. In their domain of existence, they connect Kerr BHs (that is, with no hair) with pure solitonic solutions, also known as boson stars (BS), which are regular everywhere and feature no horizons. In this sense, one can think of the KBHsSH indeed as a combined system of a BS with a horizon at its center, and therefore it shares traits of both objects.
Citation: Kleihaus et al.; 2016
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Citation: Israel 1967
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Citation: Herdeiro & Radu 2014b
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