adsabs/FOCAL huggingface dataset:
Ouchi et al. (2008) find that there is a possible excess of the Lyα LFs at z = 3.1 and 3.7 similar to the bright-end hump, and they claim that 100% of LAEs host AGNs at the bright ends of 





 and 43.4 erg s−1, respectively, based on the large-area LAE survey with the multiwavelength data set. Thus, the bright-end hump of our z = 2.2 Lyα LF may be produced by AGNs. To examine whether our LAEs at the bright end include AGNs, we use the multiwavelength data of X-ray, UV, and radio available in the SXDS, COSMOS, CDFS, HDFN, and SSA22 fields. For the X-ray data, we use the XMM-Newton source catalog in the SXDS field (Ueda et al. 2008), the Chandra1.8 Ms catalog in the COSMOS field (Elvis et al. 2009), the Chandra 4 Ms source catalog in the CDFS field (Xue et al. 2011), and the Chandra 2 Ms catalog in the HDFN field (Alexander et al. 2003). The typical sensitivity limits of these X-ray data are ∼10−16 to 10−15 erg cm −2 s−1 for the SXDS and COSMOS fields and ∼10−17 to 10−16 erg cm −2 s−1 for the CDFS and HDFN fields. We use GALEX far-UV (FUV) and near-UV (NUV) images for the UV data and obtain these images from the Multimission Archive at STScI (see also Zamojski et al. 2007 for the COSMOS field). The GALEX images reach the 3σ detection limit of ∼25–26 mag. The Very Large Array 1.4 GHz source catalogs of Simpson et al. (2006) (SXDS), Schinnerer et al. (2007) (COSMOS), and Miller et al. (2013) (CDFS) are used for the radio data. These radio data reach an rms noise level of ∼10 μJy beam−1. We find that a majority of our bright LAEs are detected in the multiwavelength data, and we summarize the numbers of these LAEs in Table 2. Under the column of “culled sample” in Table 2, we show the numbers of LAEs with no counterpart detection(s) in the X-ray, UV, and radio data. As shown in Table 2, the SXDS and COSMOS fields have the data that cover all of the X-ray, UV, and radio wavelengths. Moreover, the X-ray, UV, and radio data spatially cover the entire fields of SXDS and COSMOS with the similar sensitivities. We make a subsample that is composed of all 1576 LAEs found in the SXDS and COSMOS fields, and we refer to this subsample as SXDS+COSMOS/All. We then make another subsample consisting of 1538 LAEs with no multiwavelength counterpart detection(s) in the SXDS and COSMOS fields, which is dubbed SXDS+COSMOS/Culled.
Citation: Xue et al. 2011
Citation function: Uses; Uses; Uses; Uses
In both models, the particles interact with magnetohydrodynamic (MHD) Alfvén waves in the plasma. If the Doppler-shifted wave frequency is a constant multiple of the particle gyrofrequency in the particle guiding center frame, then a resonant interaction between the particle and the transverse component of the electric field of the MHD wave will occur (Dermer et al. 1996; Becker et al. 2006; Dermer & Menon 2009). The particle will experience either an accelerating or decelerating electric field in the transverse direction of motion over a fraction of the cyclotron period, resulting in an increase or decrease in energy. The accelerating or decelerating electric field causes the particle distributions to diffuse in energy, pushing particles to higher or lower energies in a diffusion pattern. This stochastic acceleration process typically causes the particle distributions to have a pronounced curvature in the energy spectrum (Schlickeiser 1984). The strength of the particle diffusion depends on the spectral index of the MHD turbulence, p. A Kolmogorov, p = 5/3, or a Kraichnan, p = 3/2, spectrum are most often used to model MHD turbulence. In this study, we restrict the spectral index of the turbulence to p = 2 to simulate hard sphere scattering between the MHD waves and the particle spectra. The stochastic acceleration timescale can be expressed as
3





where tdyn represents the dynamical timescale over the region in which turbulence generated (which may be smaller than the entire emission region), βA represents the Alfvén velocity of the plasma normalized to the speed of light and ξi represents the ratio of the magnetic field fluctuations relative to the background magnetic field, 





. The stochastic acceleration timescale is independent of particle mass and will therefore be the same for all charged-particle species (protons, electrons/positrons, pions, and muons). The diffusion term in Equation (1) describes the stochastic acceleration of particles in the quasi-linear approximation (Dermer et al. 1996). For gyro-resonant interactions to occur in the quasi-linear regime, the magnetic field fluctuations must be much smaller than the background magnetic field, 





. If the energy density in the plasma waves starts to approach the energy density of the magnetic field, then the field becomes disordered and there exists no well defined gyrofrequency. In both models, we use a ratio between the acceleration timescale and the escape timescale as an input parameter. The ratio between the acceleration and escape timescales constrain the maximum size in which turbulence is injected for stochastic acceleration to occur in the quasi-linear regime; see Section 5.
Citation: Schlickeiser 1984
Citation function: Uses
In order to allow us to compare the continuum shape with earlier analyses, we use phenomenological continuum models rather than more physically motivated models such as those by Becker & Wolff (2007)2 or Farinelli et al. (2016). As discussed by Müller et al. (2013), among others, phenomenological spectral models typically used to describe the continua of accreting neutron stars are the exponentially cutoff power law (cutoffpl), the power law with Fermi-Dirac cutoff (FDcut, Tanaka 1986), a negative-positive cutoff power law (NPEX, Mihara 1995), and a model consisting of a blackbody disk (diskbb, Mitsuda et al. 1984) and thermally comptonized continuum (nthcomp, Zdziarski et al. 1996; Życki et al. 1999). The residuals of the cutoffpl, FDcut, NPEX, and diskbb+nthcomp models are shown in Fig. 2, and the best-fit parameters are given in Table 1. The NPEX and FDcut residuals look very similar because they are driven to parameters that effectively mimic the cutoffpl solution. All tested continuum models describe the data similarly well. Due to its simplicity and in order to allow comparison with previous work (e.g., Younes et al. 2015), we used the cutoffpl model for all subsequent analysis. Photoelectric absorption in the interstellar medium is accounted for with the tbnew model (TBabs in XSPEC) with cross sections and abundances according to Verner et al. (1996) and Wilms et al. (2000), respectively. The iron fluorescence line complex can formally be described by a slightly broadened (



σ
=
0
.

23

−
0.04


+
0.05




$ \sigma=0.23^{+0.05}_{-0.04} $


 keV) Gaussian component at 6.59 ± 0.04 keV. This is most likely a blend of different ionization states that cannot be resolved with NuSTAR. The strongest fluorescence lines are often produced by neutral (6.4 keV), He-like (6.7 keV), and H-like iron (7.0 keV), and the structure seen in the data is also consistent with a set of narrow Kα lines from these ions, as well as neutral Kβ (7.1 keV) with a Kβ/Kα flux ratio of 13% (Palmeri et al. 2003). With fixed energies and widths, this approach is also statistically valid and has the same degrees of freedom as using one broad emission feature, but shows less interference with the continuum modeling because all line energies and widths are fixed and broadening is only due to the detector response. Using both approaches, slight residuals still remain at the iron K edge. These residuals are due to a combination of a gain-shift in NuSTAR energy calibration and the fact that the tbnew model only includes neutral iron.
Citation: Younes et al. 2015
Citation function: Compare/Contrast; Compare/Contrast
Observations over the years seem to firmly support the current acceleration of the Universe and therefore the possible existence of a generic cause responsible for it which we call dark energy (DE; see e.g. Riess et al. 1998; Perlmutter et al. 1999; WMAP collaboration 2013; Planck collaboration XVI 2014; Planck collaboration XIII 2016; Planck collaboration XIV 2016, and references therein). Cosmologists have worked hard to decipher the dark energy code, but we still ignore the physical nature of the DE and hence of the ultimate cause of the observed acceleration of the Universe. Such theoretical conundrum is the so-called Cosmological Constant Problem (CCP) (Weinberg 1989; Sahni & Starobinsky 2000; Padmanabhan 2003; Peebles & Ratra 2003; Copeland, Sami & Tsujikawa 2006; Solà 2013). In fact, the cosmological constant (CC), Λ, or equivalently the vacuum energy density associated to it, $\rho _\Lambda =\Lambda /(8\pi G)$ (G being Newton’s gravitational coupling), is usually regarded as the simplest possible explanation for the DE. Historically, the CC was introduced by Einstein in the gravitational field equations 101 yr ago (Einstein 1917). A positive, constant, tiny value (in particle physics units) of order $\rho _\Lambda \sim 2.7\times 10^{-47}$ GeV4 ∼ (2.3 × 10−3 eV)4 can explain the needed speed up of our cosmos according to the observations. The standard or ‘concordance’ cosmological model embodies such an assumption as a fundamental built-in principle, together with the hypothesis of dark matter (DM), and for this reason is called the ΛCDM model. Formulated in terms of the current cosmological parameters, the ΛCDM assumes that $\rho _\Lambda =$const. throughout the history of the Universe, with $\Omega _\Lambda \simeq 0.7$ and Ωm ≃ 0.3 at present. Unfortunately, no convincing theoretical explanation is provided about the measured value of $\rho _\Lambda$. At the end of the day, no fundamental theory, not even quantum field theory (QFT), can explain this value; and, what is worse, the typical prediction is preposterously large as compared to the measured value. The difficulties inherent to this concept were recognized as of the time when Y.B. Zeldovich first observed (Zeldovich 1967) that the contribution from QFT to the vacuum energy density should be of the order of ∼m4 for any quantum field of mass m, and therefore many orders of magnitude bigger than the existing upper bound on $\rho _\Lambda$ in those days.
Citation: Peebles & Ratra 2003
Citation function: Background
In Fig. 4, we plot the mass growth of each candidate DCBH halo as a function of redshift. In both the panels, we plot the mass of the halo versus the redshift. The left-hand panel contains haloes from the Normal simulation while the right-hand panel contains haloes from the Rarepeak simulation. The grey region in each panel below 106$\rm {M_{\odot }}~$ signifies the region below which the mass resolution of Renaissance becomes insufficient to confidently model haloes. Generally, we are able to track haloes below this threshold and into the grey region but below 106$\rm {M_{\odot }}~$ results should be treated with caution. The dashed blue line is the limit above which a halo must grow in order to overwhelm the impact of LW radiation, Mmin, LW (Machacek, Bryan & Abel 2001; O’Shea & Norman 2008; Crosby et al. 2013, 2016). The dashed red line is the approximate atomic cooling threshold, Matm, at which point cooling due to atomic hydrogen line emission becomes effective.5 Focusing first on the Normal region in the left-hand panel, we plot the growth rate of the three DCBH candidate haloes identified in the left-hand panel of Fig. 2. The DCBH candidate haloes are rapid growers but are not necessarily the fastest growing haloes in the Normal region. To emphasize this comparison, we also plot the growth of three rapidly growing haloes that contain stars. We select the three star-forming haloes from the final output of the Normal region but haloes at other redshifts do of course exist, which are rapidly growing and contain stars. In this case, we see that haloes with high dM/dz (i.e. the mass as a function of redshift) values can be star free or star forming and hence having a high dM/dz does not necessarily discriminate between DCBH halo candidates by itself. Rapidly growing haloes can become metal enriched through external enrichment processes. The enrichment allows the halo interior to cool and to form stars even in the presence of dynamical heating. Therefore, any semi-analytical model or subgrid prescription that uses dM/dz alone as a predictor for DCBH candidates will inevitably overestimate the number of candidates.
Citation: Machacek, Bryan & Abel 2001
Citation function: Uses
The gap and ring resemble those in the HL Tau system, recently found by the ALMA long baseline campaign (ALMA Partnership et al. 2015). Our result shows that gaps and rings in the (sub)millimeter dust continuum can exist, not only in relatively young disks (0.1–1 Myr) but also in relatively old disks (3–10 Myr). One possible mechanism for opening a gap is the gravitational interaction between a planet and the gas (e.g., Lin & Papaloizou 1979; Goldreich & Tremaine 1980; Fung et al. 2014). Such an interaction may also produce the spiral density waves recently found in optical and near-infrared scattered light imaging of dust grains in protoplanetary disks (e.g., Muto et al. 2012). According to recent theoretical analyses of gap structure around a planet (Kanagawa et al. 2015a, 2015b, 2016), the depth and width of the gap are controlled by the planetary mass, the turbulent viscosity, and the gas temperature. The shape of the gap is strongly influenced by angular momentum transfer via turbulent viscosity and/or instability caused by a steep pressure gradient at the edges of a gap. The observed gap has an apparent width and depth of 





 au and 





, respectively. This is too shallow and too wide compared with that predicted by theory. However, the observations are limited to an angular resolution of ∼15 au, and the depth and width could be deeper and narrower in reality. For instance, if we assume that the gap depth times the gap width retains the value derived from the observations, it is possible for the gap to have a width and depth of 





 6 au and 





, which is similar to the GPI result (Rapson et al. 2015). Such a gap could be opened by a super-Neptune-mass planet, depending on the parameters of the disk, such as the turbulent viscosity (Kanagawa et al. 2015a, 2015b, 2016). If the gap in the larger dust grains is deeper than that in the gas, the planet could be lighter than super-Neptune mass. We note that a planet of even a few Earth masses, although it cannot open a gap in the gas, can open a gap in the dust distribution if a certain amount of pebble-sized particles, whose motions are not perfectly coupled to that of gas, are scattered by the planet and/or the spiral density waves excited by the planet (Paardekooper & Mellema 2006; Muto & Inutsuka 2009).
Citation: Kanagawa et al.; 2015b
Citation function: Background; Background
Citation: Kanagawa et al.; 2015b
Citation function: Compare/Contrast
Morphology, colors, and star formation rate (SFR) primarily depend on the small-scale (1 Mpc) environment (Hogg et al. 2004; Kauffmann et al. 2004; Wetzel et al. 2012). This result has been extended by studying galaxy group samples; they show that colors and SF history most directly depend on the properties of the host dark matter (DM) halo (Blanton & Berlind 2007; Wei et al. 2010; Tinker et al. 2012), in agreement with the results of smooth particle hydrodynamic (SPH) simulations by Mazzei & Curir (2003, hereafter MC03). In this context, the investigation of the evolution of group members in the nearby universe acquires a great cosmological interest because more than half of galaxies reside in such environments. Furthermore, since the velocity dispersion of galaxies is significantly lower in groups than in clusters, the merger probability and the effects of interaction on galaxy evolution are much higher. Consequently, groups provide a zoom-in on phenomena driving the galaxy evolution before the galaxies fall into denser environments (e.g., Wilman et al. 2009; Just et al. 2010). Starting from the pioneering works of Toomre & Toomre (1972, and references therein), several papers contributed to shed light on the important role of mergers/interactions in galaxy evolution—Toomre (1977), Combes et al. (1990), Mihos & Hernquist (1994, 1996), Barnes & Hernquist (1996), Naab & Burkert (2003), Bournaud et al. (2005), and Di Matteo et al. (2007), to name a few—up to the most recent papers of Eliche-Moral et al. (2018) and Martin et al. (2018, and references therein). The dissipative merger simulations of Eliche-Moral et al. (2018) start from systems just formed, composed of a spherical nonrotating DM halo, and by a disk of gas particles with or without the presence of a stellar bulge. These simulations explore about 3–3.5 Gyr of evolution. Martin et al. (2018) focused on processes triggering galaxy transformations of massive galaxies (M > 1010 M⊙) exploiting cosmological hydrodynamic simulations by Kaviraj et al. (2017). These simulations, based on an adaptive mesh refinement code (RAMSES) and including the baryon treatment with stellar and active galactic nucleus (AGN) feedback, are able to resolve baryonic physics on kiloparsec scales, larger than we use in this paper (Section 3). They derived important statistical assessments about the processes that drive morphological transformation across cosmic time.
Citation: Wilman et al. 2009
Citation function: Motivation
Where 





 is the straggle or standard deviation of the implantation distribution, D is the diffusion coefficient, and t is the irradiation time. In the case of gas implantation via ion irradiation, it is often assumed that irradiation damage produces enough defects to serve as trapping sites that the diffusion coefficient is negligible (Martynenko 1977; Scherzer 1983). In Section 2.5, an estimate of the effective diffusion in the presence of damage produced by relativistic light ion impacts is made. Lattice defects caused by ion irradiation also provide nucleation for fixed helium bubbles (Kornelsen 1972). We can assume additional trapping due to the interaction of hydrogen with helium in the material, since helium bubbles serve as trapping sites for hydrogen atoms through synergistic effects (Hayward & Deo 2012), leading to a lower critical dose in the case of mixed hydrogen and helium exposure (Guseva & Martynenko 1981). At temperatures well below the melting point, such as that expected for a relativistic interstellar spacecraft, high-flux hydrogen irradiation damage effects have been reproduced in low-flux experiments (Gao et al. 2019). However, for an interstellar probe, the flux is sufficiently low that diffusion effects may play a non-negligible role. Classical diffusion of individual gas atoms could lead to lower local gas atom concentrations as the implantation distributions widen, or increased gas atom concentrations around trapping sites such as grain boundaries; the effect of diffusion is discussed in Section 2.5. Helium bubble nucleation and migration could lead to increased gas atom concentrations and damage near surfaces, grain boundaries, and other defects (Nakamura et al. 1977; Goodhew 1983; Lane & Goodhew 1983). No single theoretical framework exists to summarily treat these effects. However, simple models such as that presented here immediately offer compelling mitigation strategies. To perform this analysis, we find implantation profiles of ISM gas atoms at relativistic speeds using an ion–material interactions code, calculate critical concentrations for blistering onset for hydrogen and helium individually assuming a worst-case scenario of negligible diffusion, and show the effect of non-negligible diffusion on local gas concentrations.
Citation: Guseva & Martynenko 1981
Citation function: Motivation; Motivation
It is important to note that the above procedure differs from the observational one, whereby Voigt profiles are fit to the absorption features in transmission spectra (e−τ) in order to extract column densities and Doppler broadening parameters. This approach enables deblending of multiple-component absorption, and takes into account the broadening of the lines due to the instrumental profile, enabling accurate recovery of their column densities. The technique was originally devised under the premise that the intervening absorption lines in QSO spectra arose from discrete absorbing clouds, but this picture was challenged early by the smoothly distributed IGM captured in cosmological hydrodynamical simulations (e.g. Cen et al. 1994; Hernquist et al. 1996; Miralda-Escudé et al. 1996; Theuns et al. 1998; Davé et al. 1999). Absorbers that have a large spatial extent take part in the Hubble expansion, which leads to a line profile that deviates from a Voigt profile. Voigt profile fitting the transmission spectra will therefore glean slightly different results to simply taking peaks in the τ distribution. In particular, we might expect a larger number of absorbers, and some differences in the derived Doppler broadening parameters and column densities. As a check, we derived Voigt profiles for lines recovered in some of the sight-lines extracted from the simulation using our τ peak method, and plotted these on top of the transmission spectrum. An example is shown in Fig. 3. The transmission spectrum is shown as the black dashed line, and the predicted Voigt profiles are shown in green. More structure is apparent in the real spectrum, which would yield a larger number of Voigt components in Voigt profile fitting, and column densities and Doppler broadening parameters that differ slightly from those we have recovered via the τ peak method. This is an obvious caveat to our approach, which we bear in mind when interpreting our results later on.
Citation: Miralda-Escudé et al. 1996
Citation function: Background
Although wind observations are very common in AGN (see Elvis 2000; Veilleux et al. 2005; and Fabian 2012, for reviews), most studies concern ionised gas and uncertain spatial scales. In the past few years the situation changed drastically. Several fast (vOF of the order of 1000 km s-1), massive outflows of ionised, neutral and molecular gas, extended on kpc scales, have been discovered thanks to three techniques: 1) deep optical/NIR spectroscopy, mainly from integral field observations (IFU, e.g. Nesvadba et al. 2006, 2008; Alexander et al. 2010; Rupke & Veilleux 2011; Riffel & Storchi-Bergmann 2011; Cano-Diaz et al. 2012; Greene et al. 2012; Harrison et al. 2012, 2014; Liu et al. 2013a,b; Cimatti et al. 2013; Tadhunter et al. 2014; Genzel et al. 2014; Brusa et al. 2015a; Cresci et al. 2015; Carniani et al. 2015; Perna et al. 2015a,b; Zakamska et al. 2016); 2) interferometric observations in the (sub)millmetre domain (e.g. Feruglio et al. 2010, 2013a,b, 2015; Alatalo et al. 2011; Aalto et al. 2012; Cicone et al. 2012, 2014, 2015; Maiolino et al. 2012, Krips et al. 2011; Morganti et al. 2013a,b; Combes et al. 2013; Garcia-Burillo et al. 2014); and 3) far-infrared spectroscopy from Herschel (e.g. Fischer et al. 2010; Sturm et al. 2011; Veilleux et al. 2013; Spoon et al. 2013; Stone et al. 2016; Gonzalez-Alfonso et al. 2017). In addition, AGN-driven winds from the accretion disk scale up to the dusty torus are now detected routinely both in the local and in the distant Universe, as blue-shifted absorption lines in the X-ray spectra of a substantial fraction of AGN (e.g. Piconcelli et al. 2005; Kaastra et al. 2014). The most powerful of these winds, observed in 20–40% of local AGN (e.g. Tombesi et al. 2010) and in a handful of higher redshift objects (e.g. Chartas et al. 2009; Lanzuisi et al. 2012), have extreme velocities (ultra-fast outflows, UFOs, v 0.1-0.3c) and are made by highly ionised gas which can be detected only at X-ray energies.
Citation: Feruglio et al.; 2015
Citation function: Background