adsabs/WIESP2022-NER huggingface dataset:
Paper: 2017MNRAS.468.4691D (fulltext)
late He ii reionization). Therefore, they should not yet be interpreted as evidence in favour of the abundant high- z AGN models considered here. Combined with direct searches to better characterize the high- z faint AGN population, future measurements of the Ly forest opacity and of the IGM temperature will more tightly constrain the contribution of AGNs to the ionizing background, as well as their contribution to reionization. APPENDIX A: MODEL UNCERTAINTIES In this section, we discuss caveats for our models of the Ly forest opacity fluctuations considered in Section 3 . Below we argue that our models do not account for a number of effects that would, in combination, likely lead to a net decrease in the amplitude of eff fluctuations. We do not include the effects of inhomogeneous heating from patchy H i reionization. These temperature inhomogeneities would suppress fluctuations in eff because the overdense regions that were reionized earlier would tend to be cooler than average by z 5.5 (i.e. after reionization), with enhanced equilibrium H i densities from the T 0.7 dependence of the recombination coefficient. However, the ionizing background in those regions would tend to be stronger owing to the bias of the sources. The competition between these two effects would smooth out fluctuations in eff (recall that ). If H i reionization ends at z 6, this competition could have a significant effect on the width of the eff distribution (DAloisio etal. 2015 ). On the other hand, we also do not account for inhomogeneous heating during He ii reionization, which would likely work in the direction of enhancing the eff fluctuations, since the regions of strong ionizing background surrounding AGNs would also be the most recently photoheated. We do not take into account the finite lifetimes of AGNs. With a duty cycle of unity, our models effectively assume an infinite AGN lifetime. We argue that modelling the finite lifetimes could act to suppress eff fluctuations in at least two ways. First, for a fixed AGN abundance (which is set by the assumed luminosity function and magnitude limit), a smaller duty cycle would require that AGNs are hosted in haloes that are less massive and therefore less biased than in our models. This would lead to the AGN being less clustered overall, which would yield somewhat smaller eff fluctuations. Secondly, consider a region with an overdensity of bright AGNs in our fiducial model (see e.g. the lower left-hand corners of the panels in Fig. 3 ). These regions produce large segments of the Ly forest with low optical depth. Suppose now that only a fraction of these AGNs are active at a given redshift. In this case, the Ly opacity in this large-scale region would be less coherent an effect that would make the eff distribution narrower. On the other hand, we note that the proximity zone remnants left by AGNs that recently became dormant might mitigate these effects. We do not account for the anisotropy of emissions from AGNs. Our
Paper: 2019ApJ...885....4F (fulltext)
m e / m i so that the ion distribution is approximately a delta-function around p ∥ i = γ CC m i c carrying current at speed c . Details of two-stream instability are discussed in Appendix B.2 . 3.3. Ion Acceleration in the Nose We now consider the particle acceleration at the end of jets (i.e., the nose region where the jet magnetic field and current turn to the radial direction). Internal kink instability due to j z occurring in the Central Column, as analyzed in Colgate et al. ( 2015 ) and used in Section 3.2 , should persist for some distance into the nose driven by j r as the current turns radially in the nose, now less affected by curvature radiation as B φ falls like 1/ r giving curvature parameter α C ∝ 1/ r . But a new acceleration mechanism arises due to conditions in the nose region different from those in the cylindrical Central Column. One new condition concerns the sensitivity of the curvature radiation, Equation 11 (a), to R C that allows ions to be accelerated to the full voltage in the nose region, where R C = r becomes large. Another new condition concerns the escape of ions from the nose, escape itself turning out to be a necessary ingredient for the powerful acceleration mechanism not present in the Central Column. There are two ways that ions can escape from the nose. First, while ion orbits are magnetically confined in the Central Column, in the absence of an electric field ion orbits in the nose are not confined due to drifts of guiding centers in z caused by the curvature of field lines and gradients in the magnetic field intensity. Opposite drifts by electrons establish an electric charge giving E ∝ T e , the electron temperature, which can cause ions to escape by E × B motion. In addition, ions can also escape by hyper-resistive diffusion, which is what we will mainly discuss. 3.3.1. Hyper-resistive Diffusion in the Nose As in Section 3.2.1 , we represent hyper-resistivity in the nose as diffusion, noting that an electric field larger than that produced by charge separation is created if turbulence in the nose causes ions to escape by diffusion. To see this, consider the radial component of the relativistic Ohm’s law in Equations 5 (a). In the nose, with unit vector , and electron flow v e is either that due to drifts in the z -direction or that parallel to B giving so that we can drop the relativistic correction to E r . Applying Equation 5 (a) in the r -direction (perpendicular to B ) gives Equation 13 (b) giving the cosmic-ray accelerator is obtained as follows. Early in jet propagation when in the jet and nose, all of the voltage drop in the jet (giving net zero voltage around the loop) can be approximated as , which appears as an inductive electric field in our disk-centered reference frame. It is this inductive field that serves to buildup B φ behind the nose advancing at velocity . Even as the jet slows down, it continues to be true that it is that builds up B φ , giving Equation 13 (b) with in our reference
Paper: 2016MNRAS.458.2762G (fulltext)
may marginally agree with observations within errors. Figure 10. The SFR/ M s – z relation for S-I (left-hand panel), and for S-II to S-V (right-hand panel). The solid, dashed, and dotted lines show the median trends for BGGs in observations and in the SAMs of DLB07 and G11 , respectively. The grey point shows the BGG sSFR versus redshift in observations. In Fig. 11 , we show the median sSFR dependence of BGGs to the total mass of their host groups in observations (black points) and in the SAMs of G11 (red diamonds) and DLB07 (green squares), respectively. For S-I (left-hand panel), the median sSFR rapidly decreases as a function of increasing halo mass (the best-fitting linear relations corresponds to: SFR/ M s = (−4.4 ± 1.0) × M 200 + (3.8 ± 0.8). The observed trend is steeper than the model predictions. In agreement with the SAM predictions, we find no relationship between sSFR and M 200 for S-II to S-V over 0.1 z 1.3, as shown in the right-hand panel in Fig. 11 . These results suggest that, in the group scale haloes, sSFR of BGGs is a function of redshift with no halo mass dependence. Figure 11. The SFR/ M s – M 200 relation for S-I (left-hand panel), and for S-II to S-V (right-hand panel). The solid, dashed, and dotted lines show the median trends for BGGs in observations and in the SAM of DLB07 and G11 , respectively. The grey point shows the BGG sSFR versus redshift in observations. Only, for BGGs within S-I we find that the sSFR decreases as a function of halo mass. 5 SUMMARY AND CONCLUSIONS We select the BGGs in a well-defined X-ray-selected sample of galaxy groups from the COSMOS, XMM –LSS, and AEGIS fields. This sample covers a redshift range of 0.04 z 1.3 and contains about 407 galaxy groups with a mass ranging from 10 12.8 to 10 14 M ⊙ , the largest sample for such studies in the X-ray galaxy group regime. The sample allows us to study the distribution and evolution of the stellar mass and SFR of BGGs and the relation among the stellar mass, SFR, and halo mass. We compare our results with sample drawn from the Millennium simulations with three SAMs ( G11 , DLB07 , and B06 ). We summarize our results as follows. The shape of the stellar mass distribution of the BGGs evolves towards a normal distribution as the groups evolve. This evolution is quantified by the skewness and Kurtosis of the stellar mass distribution and their redshift dependence. At all redshifts, we find that the stellar mass distribution tends to skew towards low masses. The observed fraction of the BGGs shows a strong deviation from the Gaussian fit at ≈10 10.2 M ⊙ , which are generally SF/young population. In contrast, the SAM predictions show no redshift evolution of the stellar mass distribution. Within the probed SAMs, the shape of the stellar mass distribution predicted by B06 model is more consistent with the observations at z 0.5. We show that the average stellar mass of the BGGs evolves with redshift by a factor of ∼2 from z = 1.3 to the present day. At z
Paper: 2016MNRAS.460.2286S (fulltext)
This is a great advantage of such a Nasmyth system. Any optical designer of telescopes can observe that in this Nasmyth system the exit pupil is near the plane fold mirror, and that if some small adjustments are made to this system, an AO deformable mirror can be added. After this AO deformable mirror, a plane fold mirror can still be added to reflect light to the Nasmyth focus. In such an arrangement the length of path from the exit pupil to the pupil formed by the secondary mirror and that from the exit pupil to the Nasmyth focus should be equal. These two plane mirrors can be used, respectively, as AO deformable mirror and tip-tilt mirror for all instruments at Nasmyth platform and coudé focus. This is the second great advantage of this Nasmyth system. Such a system has been applied in the Nasmyth system of 39 m European Extremely Large Telescope (Delabre 2008 ; ESO 2011 ). However, if an AO deformable mirror is added in such a Nasmyth system, there will be some shortcomings as well. The size of the AO deformable mirror is bigger than those AO deformable mirrors set up separately in Nasmyth instruments and coudé instruments. It is difficult for a big AO deformation mirror to perform the high temporal frequency movement. Besides, an AO mirror is unnecessary for some Nasmyth instruments and coudé instruments. As it is located in the tube of the telescope, this AO mirror will produce some harmful thermal effects, especially since this mirror needs more power due to its large size. Besides, servicing also possess some difficulty. Given these circumstances, at present we only intend to use one plane fold mirror, which will serve as a tip-tilt mirror in our Nasmyth system. Nevertheless we insist that the addition of an AO deformable mirror in this Nasmyth system is an important advantage and we will continue to consider whether or not to make such an addition during future research and design of Nasmyth and coudé instruments. Excellent image quality has been obtained in our Nasmyth system. In the next section one can witness the excellent image quality also in our coudé system design. We are convinced that when we make some small readjustments for adding an AO deformable mirror, the excellent image quality can still be obtained in Nasmyth and coudé systems. In Nasmyth platforms many instruments will be installed. In order to meet this challenge, we propose (Su, Wang Cui 2004b ) the use of four Nasmyth platforms in this telescope, i.e. two Nasmyth platforms I on the upper level and two Nasmyth platforms II on the lower level. In such an option, the two light bundles to Nasmyth platform II come perpendicularly up (or down) towards this platform; for some instruments that is an advantage. Since the lengths of paths from secondary mirror to Nasmyth I foci and to Nasmyth II foci may be different, M2 and M3 will be moved to make the image arrive at the required foci. In these situations fairly large off-axis aberration may occur and FOV will be
Paper: 2020A&A...642A.132C (fulltext)
acted at different temporal and spatial scales throughout the lifetime of the galaxies. The size of our sample does not allow for a separate analysis of whether AGNs or bars can play a role in the observed stellar metallicity. However, we observe that the mere presence of an AGN at z = 0 does not necessarily imply that it has played a role in shaping the observed galaxy stellar metallicity. An AGN can be a transient phenomenon in a galaxy. With AGN duty cycles spanning timescales in the range 10 6 − 10 8 yr ( Haehnelt Rees 1993 ; Davis et al. 2014 ; Storchi-Bergmann Schnorr-Müller 2019 ), we expect that the AGN effects on stellar metallicity should be due to past and not current AGN activity. Regarding bars, there is no consensus in the literature as to the longevity of bars, nor how many bar events the average galaxy experiences during its lifetime ( Sellwood 1999 ; Athanassoula 2002 ; Bournaud Combes 2002 ; Elmegreen et al. 2004 ; Combes 2004 ; Regan Teuben 2004 ; Pérez et al. 2008 ; James Percival 2016 ). While several authors have observed a correlation between gas abundance gradients and the presence of bars, in particular, a flattening of the gradient ( Vila-Costas Edmunds 1992 ; Martin Roy 1994 ; Zaritsky et al. 1994 ), more recent works have found no evidence of such a correlation by analysing both gas-phase metallicity and stellar metallicity ( Sánchez et al. 2012b , 2014 ; Sánchez-Blázquez et al. 2014 ; Cheung et al. 2015 ; Sánchez-Menguiano et al. 2016 ). The upper panels of Fig. 2 show the light-weighted metallicity as a function of stellar mass surface density log(Σ ⋆ ) for galaxies in the field and in groups. Field galaxies shown in Fig. 2 are a random realisation of 107 galaxies from the field sample, which were selected to have the same overall mass distribution as the group sample. Each galaxy, in either of the two samples, contributes several points to this figure. The number of points varies from galaxy to galaxy, as does the number of radial bins that each profile has. Points are colour-coded in tones of grey according to their radial distance in terms of the effective radius, r / r e . The solid curves show the number-density iso-contour levels. The stellar mass surface density and metallicity are strongly correlated although we observe significant scatter. To compare galaxies in the field and groups, the upper right panel of Fig. 2 shows the iso-contour for each environment considered. We clearly observe that galaxies in groups are more metallic for a fixed value of stellar mass surface density. Fig. 2. Top panels : light-weighted metallicity as a function of stellar mass surface density Σ ⋆ . Dots are colour-coded in tones of grey according to their distance r / r e . The left panel corresponds to galaxies in the field, the central panel shows galaxies in groups, and the right panel shows the iso-contour levels for galaxies in groups (red lines) and galaxies in the field (blue lines). The bottom panels show the median
Paper: 2018ApJ...862..157G (fulltext)
moon would be able to create the gap. These viscosities are notably lower than those measured by Tiscareno et al. ( 2007 ) and Colwell et al. ( 2009a ) in this region. For , the ratio has a value of 1.2 so the moon might just be able to create the Kuiper gap but for a slightly larger viscosity of, for example, , this ratio has a value of only 0.12 and the moon should definitely create a propeller and not a gap around the entire circumference of the ring. Let us consider the wide Bessel gap as another example. The diameters a spherical ice moon (or hypothetically even a silicate moon) would have to have to open this gap are and , respectively. Even for the lower limit of , the moon would have to have a nonsensical density of to fall below a diameter of . We therefore believe that single solid icy or silicate bodies can be ruled out as an explanation for the observed gaps in the C ring and the Cassini division. Lissauer et al. ( 1981 ) described the rift and gap structure of the Cassini division with a set of five embedded moons. They assumed a dispersion velocity with a spatially constant value of , which for , yields , which is much larger than the viscosities measured by Tiscareno et al. ( 2007 ) and Colwell et al. ( 2009a ), thus our chosen lower limit of . Consequently, the moons they predicted () are significantly larger than the ones we predict (see Table 1 ). Those moons might have been overlooked by Voyager , but would definitely have been detected by the cameras of the Cassini spacecraft. Ruling out the existence of single icy bodies as an explanation for the observed gaps in the C ring and the Cassini division raises the question what process could explain their formation. We believe that a swarm of icy collision shards much smaller than the dimensions presented in Table 1 , but still much larger than the largest ring particles could scatter the ring particles away, creating a gap that despite the existence of this swarm would appear completely void to the cameras of the Cassini spacecraft. However, this hypothesis cannot be tested easily with the model presented in this article and requires additional research. 4. Summary We developed a nonlinear diffusion model for gaps in planetary rings that accounts for gravitational scattering by the embedded moon and for the viscous diffusion of the ring material. This model can be used to estimate the ring’s viscosity from the radial density profile of the gap the moon creates knowing its radial position and mass and the power-law exponent β (). The mass of the moon can be derived by analyzing the wake region (Showalter et al. 1986 ; Seiß et al. 2010 ). We applied the model to the Encke and Keeler gap created by Pan and Daphnis to estimate the rings viscosity. Then, we applied the model to the gaps in the C ring and Cassini division to estimate the masses of moons that would be able to create those gaps. For , we estimate for the Encke gap, which is not an unexpected result, as Tiscareno et
Paper: 2020ApJ...900..125R (fulltext)
NRAO Astronomical Image Processing System 17 17 http://www.aips.nrao.edu/ package. In order to avoid false associations of radio matches, we downloaded the images from DSS1 and DSS2, which are digitized versions of several photographic astronomical surveys, in addition to radio maps, which can be found at the ESO archive. 18 18 http://archive.eso.org/dss/dss Comparison of radio and optical sources on maps from both domains, together with the NASA/IPAC Extragalactic Database, 19 19 https://ned.ipac.caltech.edu/ enabled us to distinguish incorrect matches. As some sources turned out to have extremely extended, i.e., beyond 3′, radio structures we were gradually increasing the radio search by 1′ as long as the association for the whole structure was found. Once the whole radio-emitting region was identified, the radio flux of each of the components was summed up and assigned to the given source. For sources lacking radio detections, we assigned them upper limits corresponding to the value of the sensitivity of the survey that contains the source in its footprint. While the sensitivity of NVSS and SUMSS is the same, for objects located in the area covered by both NVSS and FIRST, we decided on FIRST upper limits, as its sensitivity is lower than that of NVSS. Because the radio catalogs we used were conducted at two different frequencies, the radio fluxes at 843 MHz were recalibrated to 1.4 GHz using a radio spectral index of α r = 0.8 (with the convention of F ν ∝ ν − α ) so that the rest of our calculations are consistent. A.2. Mid-infrared Data The mid-infrared measurements were taken from the AllWISE Data Release (Cutri et al. 2013 ) which, by combining data from the cryogenic WISE (Wright et al. 2010 ) and post-cryogenic NEOWISE (“near-Earth object + WISE,” Mainzer et al. 2011 ) survey phases, resulted in a comprehensive view of the mid-infrared sky. Out of four available bands (at 3.4, 4.6, 12, and 22 μ m, corresponding to the W1, W2, W3, and W4 bands, respectively), we decided to make use of the W3 band, which was driven by the fact that at this specific wavelength the dusty torus is transparent enough to observe radiation coming directly from the central region of an AGN. At the shorter wavelengths of the W1 and W2 bands, the dusty, circumnuclear tori are optically thick and radiate anisotropically (Hönig et al. 2011 ; Netzer 2015 ). At the longer wavelengths of the W4 band, the dusty torus becomes even more transparent, but its measurements are affected by much larger errors than in W3 because (1) the W4 band traces the warm dust continuum at 22 μ m and can be contributed to by starbursts (Ichikawa et al. 2019 ); (2) the sensitivity in the W4 band is much lower than in the W3 band and its signal-to-noise ratio is the lowest of all the WISE channels, resulting in a lower detected fraction for the objects in our sample; and (3) the resolution of W4 images is worse than in other bands (e.g., 12′′ in W4 versus 65 in W3), so for some of our
Paper: 2019ApJ...873...35S (fulltext)
capabilities on large-aperture telescopes has revolutionized the study of high-redshift ( z ≳ 5.7) quasars (e.g., Fan et al. 2001 , 2003 , 2004 , 2006a ; Cool et al. 2006 ; Goto 2006 ; McGreer et al. 2006 ; Venemans et al. 2007 , 2015 ; Willott et al. 2007 , 2009 , 2010b ; Jiang et al. 2008 , 2009 , 2015 , 2016 ; Mortlock et al. 2009 , 2011 ; Morganson et al. 2012 ; Bañados et al. 2014 , 2016 ; Matsuoka et al. 2016 , 2018a , 2018b ; Reed et al. 2017 ; Tang et al. 2017 ). These earliest quasars, powered by supermassive black holes (SMBHs), not only constrain the physics of black hole (BH) accretion and host galaxy assembly at cosmic dawn but also provide critical information on the physical conditions of the intergalactic medium in the early universe (e.g., Fan et al. 2006a ). Since the first successful identification of z ≳ 6 quasars (Fan et al. 2001 ), the searches for these high- z quasars have matured and dramatically increased the inventory of these rare objects. There are now more than 250 quasars known at z > 5.6, with the most distant quasar reaching z ∼ 7.5 (Bañados et al. 2018b ). These impressive statistics now enable a transition from individual case studies to ensemble studies of high- z quasars. The vast majority of z ≳ 6 quasars are selected from wide-field optical+IR imaging with the dropout technique and are confirmed with optical spectroscopy (e.g., Fan et al. 2001 ). The confirmation optical spectroscopy provides limited information about these high- z quasars themselves, and to probe the physical properties of these systems (e.g., BH mass, spectral properties, broad and narrow absorption lines, etc.), near-IR spectroscopy is necessary to cover the rest-frame UV from C iv 1549 Å to Mg ii 2798 Å. Some earlier near-IR spectra of z ≳ 6 quasars already provided a glimpse of their physical properties, such as BH masses and Eddington ratios (e.g., Jiang et al. 2007 ; Kurk et al. 2007 ; Willott et al. 2010a ; De Rosa et al. 2014 ; Mazzucchelli et al. 2017 ; Kim et al. 2018 ) and chemical abundances in the broad-line region (e.g., Barth et al. 2003 ; Kurk et al. 2007 ; De Rosa et al. 2011 ). These measurements of BH properties also facilitate the study of the relations between BH growth and host galaxy assembly in the early universe, where the host galaxy properties are mostly inferred from the molecular gas emission in the millimeter regime (e.g., Walter et al. 2003 , 2009 ; Riechers et al. 2009 ; Wang et al. 2013 ; Venemans et al. 2016 ; Decarli et al. 2017 , 2018 ; Willott et al. 2017 ). However, near-IR spectroscopy of z ≳ 6 quasars is expensive, and thus only a small fraction of them have existing near-IR spectra, some of which are also of low quality (e.g., insufficient spectral coverage, signal-to-noise ratio, and spectral resolution). Given the importance of understanding the physical properties and growth of these high- z SMBHs, we have assembled a large sample of high-quality near-IR spectra for these objects, as