adsabs/WIESP2022-NER huggingface dataset:
Paper: 2020MNRAS.492.1991M (fulltext)
side and top histograms show that our SINFONI targets (red-hatched histogram) form a fair sub-sample of the full CARLA sample (blue-shaded histogram) in radio power, UV luminosity, and BH mass, even though the SINFONI sample covers a more limited range in redshift. A list of the targets and relevant observational information in our quasar sample is given in Table 1 . Table 2 summarizes the main parameters determined for the six broad-line radio galaxies from Nesvadba et al. ( 2011 ) that were used in our analysis. The available parameters include H α-based BH mass estimates and Eddington ratios that were determined in a similar fashion to the analysis performed on the CARLA quasars in this paper. Table 1. Sample information. Note . a Redshifts are from SDSS. Table 2. Broad-line radio galaxies from Nesvadba et al. ( 2011 ). Notes . a Redshifts based on H α taken from Nesvadba et al. ( 2011 ). † Radio luminosity at 325 MHz taken from De Breuck et al. ( 2000 ). ‡ Radio luminosity at 500 MHz taken from Wylezalek et al. ( 2013 ). 2.2 Observations We used SINFONI at UT4 of the VLT at the European Southern Observatory (ESO) to observe a sample of 35 high-redshift quasars. The purpose of the project was to obtain NIR spectra probing the rest-frame optical region around H α in order to obtain more accurate BH masses for quasars selected from the CARLA survey and SDSS. The data were taken over 6 yer (2009–2015) as part of programmes 095.B-0323(A), 094.B-0105(A), 093.B-0084(A), 092.B-0565(A), 091.B-0112(A), 090.B-0674(A), and 089.B-0433(A) (PI: Jaron Kurk). Except for 089.B-0433(A), all programmes targeted radio-loud quasars from CARLA. The 089.B-0433(A) targets were selected from SDSS on the basis of their very high (C iv -based) BH masses but are not part of the CARLA sample (see Section 2.1 ). We used the SINFONI K band (1.95–2.45 μ m) with the spatial scale of 125 × 250 mas (without adaptive optics). This setup delivers a spatial aperture of 8 × 8 arcsec 2 with a spectral resolution of ∼400 km s −1 , allowing a robust sky subtraction. The observational strategy includes on-source dithering (ABBA pattern) in order to optimize the exposure time on the source and sky subtraction, and a telluric standard star for telluric band removal and flux calibration. We took 14–16 individual exposures of 300 s each, resulting in total integration times of 4200–4800 s. 2.3 Data reduction The data reduction was performed using the SINFONI pipeline v2.9.0 (Modigliani et al. 2007 ) with the Reflex interface. The pipeline creates master calibration images and applies them to the observations, starting with a non-linearity map constructed from the flat-field frames. A master dark frame is created from the dark frames and the hot pixels are mapped. Individual flat-field images are combined into a master flat. Optical distortion and slitlet distances are computed, and the dispersion solution for the wavelength calibration is determined from arc spectra. Finally,
Paper: 2017ApJ...848...15F (fulltext)
neutrinos in spatial and temporal coincidences were reported by the IceCube neutrino telescope around this burst. A similar powerful burst GRB 130427A with energy of erg was detected by several satellites and ground-based telescopes (Ackermann et al. 2014 ; Maselli et al. 2014 ; Vestrand et al. 2014 ) and although searches for TeV–PeV neutrinos were performed, no excess were found above background. Gao et al. ( 2013b ) stated that the neutrino non-detection could constrain the values of the bulk Lorentz factor, emitting radius and the energy fraction converted into cosmic rays . They found that almost independently of the bulk Lorentz factor, the energy fraction between electrons and cosmic rays lies in the range . Although a robust analysis could be required, a simple proof can be done for GRB 160625B following a similar procedure. Form our results obtained in early afterglow can be seen that the energy fraction given to accelerate electrons and amplified the magnetic field at the end of the prompt phase is and , respectively. Taking into consideration the energy conservation condition , then the energy fraction converted into cosmic rays would be limited by . This result is very similar to that found by Gao et al. ( 2013b ) for GRB 130427A and might explain the lack of high-energy neutrinos around GRB 160625B. 7. Kann et al. ( 2010 ) studied the optical photometry data in a total of 42 GRB afterglows. They found that 10% of the afterglows presented optical peaks followed by a fast decay which are usually associated with a reverse-shock flash. Several authors have claimed that this kind of afterglow, as observed in GRB 080319B (Racusin et al. 2008 ), GRB 130427A (Ackermann et al. 2014 ; Vestrand et al. 2014 ), GRB 050904 (Kann et al. 2007 ), GRB 120711A (Martin-Carrillo et al. 2014 ), and GRB 990123 (Akerlof et al. 1999 ), among others, are only present in the most luminous bursts. Given that GRB 160625B has been one of the most powerful bursts detected which exhibited an optical flash with a fast decay, this burst seems to confirm this statement and belong in the same category. Table 2 Parameters found of the Early and Late Afterglow Table 3 Quantities Derived with our Leptonic Model for the Early and Late Afterglow Note. Quantities for the early afterglow are calculated at 150 s and for the late afterglow at . At the jet-break time ( s), the jet opening angle and the bulk Lorentz factor are and 6.9, respectively.) 5. Conclusions We have described the non-thermal multiwavelength observations of GRB 160625B collected with Fermi -LAT, Swift -XRT and several optical ground observatories. The multiwavelength observations of the early afterglow are consistent with the afterglow evolution in a stellar wind medium. The optical spectral index is consistent with the synchrotron radiation while GeV γ -ray flux with SSC emission dominated by the high latitude emission. In this event, a strong correlation between GeV γ -ray and optical fluxes
Paper: 2017MNRAS.471..422T (fulltext)
X1, the most well known of these systems, periodic occultation of the source due to the precession of the warped inner disc results in a luminosity not proportional to the accretion rate and as a consequence the torque luminosity correlation is not good (Klochkov etal. 2009 ). We have seen that the dimensionless torque near the equilibrium is qualitatively different from the predictions of the existing models. Across the torque equilibrium the torque is concave-up in its dependence on the fastness parameter, whereas the existing models (Ghosh Lamb 1979b ; Li Wang 1996 ) predict concave-down (convex) relation. The torque has a cubic dependence near the torque equilibrium and the general qualitative behaviour can be represented by superposition of two hyperbolic functions and is shown to be very different from the existing analytical models. These qualitative differences indicate either that an important ingredient of the discmagnetosphere interaction is not captured in these models or that one or more of our assumptions in constructing the torque are not justified. An assumption implicit in our construction of the torque is the proportionality of the inner radius with the Alfvn radius, R in R A , yet it is possible that is not a constant but depends on the fastness parameter. This condition will be relaxed in a future study, but we would like to mention that our comparison of the dimensionless torque with the existing torque models of Ghosh Lamb ( 1979b ) and Li Wang ( 1996 ) is self-consistent in the sense that these authors also find to be constant, independent of the fastness parameter. Another assumption that can be questioned is the proportionality of luminosity, and hence the flux received, to the accretion rate. There could be two ways this assumption may fail: it is possible that the inner disc is warped and obscures some of the X-ray flux as is the case in Her X1. This possibility is recently suggested by Beri, Paul Dewangan ( 2015 ) relying on the pulse-phase dependence of the emission lines. Yet we note that the object shows a very smooth torque evolution near the torque equilibrium with a very satisfactory torqueflux correlation (Camero-Arranz etal. 2010 , 2012 ) as opposed to the case in Her X-1. This is probably because we are seeing the system face on: we do not see the Doppler modulations, yet we do know the orbital period through optical modulation of the companion (Chakrabarty 1998 ). There is another way the linear relation between the accretion rate and the luminosity could break down. Yi, Wheeler Vishniac ( 1997 ) suggested that the torque reversals from spin-up to spin-down could be a consequence of transition of the flow in the disc from a Keplerian flow to a sub-Keplerian flow. A torque reversal from spin-up to spin-down is then associated with a spectral transition from soft- to hard-state (Vaughan Kitamoto 1997 ; Yi Vishniac 1999 ). Such spectral transitions, in the case of black holes, are associated with transition
Paper: 2019MNRAS.486.3423H (fulltext)
external stream, they nevertheless dominate the broad-band opacity at FUV, optical, and IR wavelengths if they are present. They also dominate at EUV wavelengths (), if the hydrogen neutral fraction is less than ≈0.001. The strong coupling assumption means that all the radiative forces applied to the dust grains are directly felt by the gas also. 2.1 Bow shells supported by radiation and wind The internal pressure is the sum of wind ram pressure and the effective radiation pressure that acts on the bow shell. The radiative momentum loss rate of the star is L / c and the wind momentum loss rate can be expressed as where η w is the momentum efficiency of the wind, which is 1 in all cases except WR stars (Lamers Cassinelli 1999 ). If the optical depth at UV/optical wavelengths is sufficiently large, then all of the stellar radiative momentum, emitted with rate L / c , is trapped by the bow shell. At the infrared wavelengths where the absorbed radiation is re-emitted, the grain opacity is much lower, so the bow shell is optically thin to this radiation. Combined with the fact that the UV single-scattering albedo is only ≈0.5, this means that the single-scattering limit is approximately valid, and we adopt it here. We define a fiducial bow shock radius R * in the optically thick, radiation-only limit by balancing stellar radiation pressure against stream ram pressure at the bow apex along the symmetry axis from the star: which yields We now consider the opposite, optically thin limit. If the total opacity (gas plus dust) per total mass (gas plus dust) is κ (with units of ), then the radiative acceleration is Therefore, an incoming stream with initial velocity, v ∞ , can be brought to rest by radiation alone at a distance R ** where yielding On the other hand, we can also argue as in the optically thick case above by approximating the bow shell as a surface, and balancing stellar radiation pressure against the ram pressure of the incoming stream. The important difference when the shell is not optically thick is that only a fraction 1 − e −τ of the radiative momentum is absorbed by the bow, so that equation ( 2 ) is replaced with In the optically thin limit, 1 − e −τ ≈ τ, so these two descriptions can be seen to agree ( R 0 → R ** ) so long as which we will assume to hold generally. We define a fiducial optical depth, With these fiducial values established, we return to the combined radiation plus wind scenario. Adding the stellar wind ram pressure term from equation ( 1 ) allows us to write the general bow shell radius in terms of the fiducial radius as where x is the solution of Since this is a transcendental equation, x must be found numerically, but we can write explicit expressions for three limiting cases: The first case, x ≈ (1 + η w ) 1/2 , corresponds to a radiation bow shock (RBS); the second case, x ≈ 2τ * , corresponds to a radiation bow wave (RBW); and the third case, , corresponds to a wind bow shock (WBS). The two bow shock cases are
Paper: 2017ApJ...839..120W (fulltext)
the host galaxy potential. The average difference between AGNs and SF galaxies is 0.16 dex. Note that even among the objects with single Gaussian [O iii ], we clearly detect the systematic difference of the ratio between AGNs and SF galaxies (the average difference is 0.11 dex). 3.2. [O <sc>iii</sc>] Velocity Dispersion versus Luminosity We investigate how the [O iii ] kinematics are related to [O iii ] luminosity for AGNs and SF galaxies in Figure 5 . There is an increasing trend of σ O iii with increasing luminosity in AGNs, though the trend is very broad. The increasing trend is also present among composite objects while the the luminosity range of [O iii ] is an order of magnitude narrower. These trends show that gas outflows are stronger in AGNs with higher luminosity, indicating that AGN energetics are responsible for the outflow kinematics. In contrast, we find no clear increasing trend among SF galaxies. On average σ O iii is lower in SF galaxies, while the range of σ O iii is as broad as that of AGNs. Since the average stellar mass of AGNs and SF galaxies is very different, as we pointed out in S.2.2, we plot the same comparison after limiting stellar mass between 10 10 and 10 10.5 (bottom panels in Figure 5 ). Again, we find that AGNs and SF galaxies show different trends, suggesting that these different trends are not due to the difference in mass scales between AGN and SF galaxy samples. Figure 5. Comparison of [O iii ] velocity dispersion with [O iii ] luminosity, respectively, for AGNs (left), composite objects (middle), and SF galaxies (right), using the entire sample (top) and the stellar mass-limited sample (bottom). The colors represent the number density in each bin with Δ x = 0.03 dex and Δ y = 0.01 dex. 3.3. [O <sc>iii</sc>] Velocity–Velocity Dispersion Diagram In Figure 6 , we combine the velocity and velocity dispersion of [O iii ] to investigate the difference of gas kinematics between AGNs and SF galaxies. We find a dramatic contrast between AGNs and SF galaxies in terms of the range of gas velocities. While SF galaxies are concentrated at the lower center with a limited range of velocity shift and velocity dispersion ( σ O iii ∼150 km s −1 ), AGNs show a much larger distribution, including AGNs with extreme kinematics (i.e., and σ O iii > 400 km s −1 ). Figure 6. [O iii ] Velocity–velocity dispersion diagrams, respectively, for AGNs (left), composite objects (middle), and SF galaxies (right), using the total sample (top) and the stellar mass-limited subsample (bottom). The average σ O iii is 175.1 ± 76, 125.3 ± 49.9, and 73.2 ± 28.6 km s −1 , respectively, for AGNs, composite objects, and SF galaxies, indicating that AGNs have more than a factor of two broader [O iii ] than SF galaxies. Also, the dispersion of the σ O iii distribution is significantly different with a factor of ∼3 broader distribution in AGNs. Note that the σ O iii distribution is skewed toward high values, indicating the presence of extreme
Paper: 2018A&A...619A.144P (fulltext)
in determining the fractional abundance of atomic hydrogen in dark clouds. The main results are listed below. (i) The CR dissociation rate, ζ diss , is primarily determined by secondary electrons produced during the primary CR ionisation process. These secondary electrons can efficiently dissociate H 2 and represent the only source of atomic hydrogen at column densities higher than ≈10 21 cm −2 , regulating the ratio in dark clouds. (ii) ζ diss entering the balance Eq. ( 4 ) is not equal to the ionisation rate ζ ion , as assumed in some previous work. We find that the ratio ζ diss ∕ ζ ion varies between ≈0.63 and ≈0.7, depending on the column density range, while ζ ion is a decreasing function of the column density. (iii) Even given the uncertainties in the values of H 2 formation rate, temperature, total hydrogen volume density, and IS UV radiation field for each cloud, only a CR proton spectrum enhanced at low energies (such as our model ) is capable of reproducing the upper values of measured f H under the assumption of steady state. We note that neither model nor model is able toreproduce the entire set of observational data: the spread in the values of f H at any given column density must be attributed to time dependence or to individual characteristics of each cloud. For example, tangled magnetic field lines and/or higher volume densities would result in a stronger CR attenuation and thereforein a lower f H . (iv) An accurate description of H 2 dissociation in dense environments is essential because many chemical processes (such as CO hydrogenation and its depletion degree onto dust grains, or formation of complex organic molecules) critically depend on the abundance of atomic hydrogen. Appendix A Equilibrium distribution of protons and H<sub>fast</sub> atoms at low energies Because of the process of electron capture at low energies, CR protons interacting with H 2 are efficiently neutralised, creating fast H atoms. In parentheses we list the cross section of the respective process. At the same time, H fast atoms reacting with H 2 yield The electron capture by protons, reaction ( A.1 ), and the reaction of H fast self-ionisation, reaction ( A.3 ), are catastrophic processes, since the respective projectile particles disappear after such collisions. The reactions of H 2 ionisation and dissociation by H fast atoms, reactions ( A.2 ) and ( A.4 ), respectively, are continuous loss processes, where the projectile kinetic energy decreases only slightly after each collision. The efficiency of continuous energy losses is generally characterised by the projectile stopping range (see, e.g. Padovani et al. 2009 ). For our calculations, is taken from Rudd et al. (1983) , is adopted from Phelps (1990) and Kunc Soon (1991) , is computed by Stier Barnett (1956) , van Zyl et al. (1981) , and Phelps (1990) , and is taken from Dove Mandy (1986) and Esposito Capitelli (2009) . In Fig. A.1 , we plot the cross sections and the inverse of the proton
Paper: 2016ApJ...832....7A (fulltext)
times. No physical prescriptions govern these behaviors, only the universe’s observed SFH and the end-states of a set of real objects. In this sense, the G13 “model” is not a model at all, but a realization of possible lognormal SFHs for a sample of galaxies. As such, all physics must be inferred from the resulting distribution, or read into the form of the SFH. Sections 3.2 , 3.3 , and 4 revisit this point. Note: Given its mass limit, G13 does not describe the SFHs of local dwarfs (though neither does it contain any). Unlike giant systems (Rodighiero et al. 2011 ), these are known to grow via stochastic bursts (Weisz et al. 2014 ) which a single lognormal cannot capture. Hence, the following discussion is likely not applicable to them. Our principal interest lies in understanding what shapes the longer-term growth modes of galaxies that have consistently dominated the universe’s mass and budget (Section 2.1 ). 2.3. Previous Descriptive Successes In G13 , we showed that the model naturally reproduced the population of galaxies O13 (and subsequently Dressler et al. 2016 ) found to have both late-peaking and narrow SFHs, something impossible to do using τ or delayed- τ SFH parameterizations (Figures 2 and 9 ; O13 ). This encouraged us that the G13 model represented not only an improved description, but perhaps a physically meaningful worldview. More encouragement has come as the model continues to reproduce a range of observations it was neither constrained by nor intended to describe. We explore most of these results below, but review here what we have demonstrated to date: 1. Constrained only by the cosmic SFRD and input galaxies, G13 (Figure 7 ) showed that the model faithfully reproduced distributions at four epochs at . As O13 did for τ /delayed- τ models, G13 showed that Gaussian SFHs failed this test badly. 2. Turning the distributions into modeling constraints and refitting for each galaxy’s , G13 (Figure 13 ) showed that the model reproduced the – relation (the “Star Formation Main Sequence,” or “SFMS”) from Daddi et al. ( 2007 ). 9 9 The agreement with distributions achieved by the initial model suggests that the SFMS could have been reproduced without turning these into constraints. We have not performed this test, however. 3. Similarly constrained at , Abramson et al. ( 2015 ) showed that the model reproduced the evolution of the stellar mass functions of star-forming and non-star-forming galaxies at . This was despite the fact that no mass-sensitive constraints were ever imposed except via the intrinsic mass distribution of the input data. These achievements by no means make the G13 model uniquely successful: others are equally capable (e.g., Sparre et al. 2015 ). However , as G13 entails no explicit treatment of any physical process, we are spurred to investigate what makes this simple mathematical description an accurate expression of complex, poorly understood underlying physics. 2.4. New Challenges Here, we test the G13
Paper: 2019MNRAS.486.1805Z (fulltext)
the inner initial region of the debris disc is chosen as to coincide with the appropriate tidal radius given the NS and WD properties. However, the exact structure of the debris disc is more complex than our simplified models and we therefore checked the dependence of slightly by moving the disc position inwards or outwards (see R 0 / r t in Table 2 ). As we show, the results depend on the chosen inner positions, but the overall behaviour is robust and is not strongly dependent on the exact choice. Closer-in debris discs are more effective in allowing for material to accrete into close-in orbits before outflows become efficient enough to quench further accretion. The overall ejected mass and total energetic are therefore determined by both the overall NS and WD mass and the inner radius of the disc; these are well reflected in Table 2 . Overall the nuclear energy deposited increases with increasing density and with Helium content. These provide more favourable conditions for nuclear burning. The density is determined by the total mass of the debris disc (the disrupted WD), the inner radius ( r d ), the disc scale height, and to some extent the disc composition that affects the EOS (see equation 3 ). As can be seen in Table 2 , the nuclear energy produced follows these various trends; it increases with increasing WD mass and Helium content, as well as with decreasing disc scale height and inner radius. The tidal radius, which determines the inner radius, depends on both the NS mass and the WD mass, as well as on the WD composition, and hence more massive NSs (giving rise to larger tidal radii) correspond to lesser production of nuclear energy. The pure Helium WD case (not shown) is the only case where no nuclear burning is observed. This may appear counter-intuitive given the lower temperatures and densities required to ignite Helium, however such WDs have a lower mass and are much puffier, hence the effective density of the debris disc they produce is far lower than that of CO WD debris disc, explaining the result. As discussed above, our results suggest that only a small fraction of the WD debris disc is effectively ejected during the modelled evolution, while most of the material expands to produce a more isotropic configuration, but stays bound to the NS. The long-term evolution of such expanded debris ‘atmosphere’ is beyond the scope (and computational capabilities) of the current study focusing on the immediate outcomes of the merger following the disruption of the WD. As already noted, nuclear burning is not the main driver of the outflows in NS–WD mergers, and only a small fraction of the WD debris disc experiences any nuclear burning. Hence, the majority of the ejected unbound material (identified through the comparison of the kinetic and potential energy of each mass element) shows the same composition as the original WD, i.e. C/O composition (and He in the hybrid WD cases). However, most of the unbound material is ejected through