adsabs/WIESP2022-NER huggingface dataset:
Paper: 2016ApJ...832...74Z (fulltext)
between the curves of Figure 6 (a) are due to this anisotropic and coherent velocity strength. The Fourier transform as a function of k y (not shown) is almost the same as Figure 6 (a). Figure 6 (c) shows the kinetic energy power spectrum as a function of ; the differences between the energy contribution of the modes in the cases with and without shear are more pronounced, but still rather small. We speculate that the action of the stress tensor (when the shear is applied) leads to an increase of mode contributions at all scales, but in particular near the diffusion region, increasing the reconnection rate; this is in agreement with Nakamura et al. ( 2008 ), who reported an enhancement of island production when shear was present. Figure 6. (a) Power spectra of the velocity as a function of k . (b) Velocity as a function of k x . (c) Kinetic energy as a function of k . For all cases the red line represents the fan stage at t = 303.0 s just before the shear is applied, and the blue one just after it, at t = 304.0 s. If we consider a magnetic configuration scheme (see Figure 1 of Paper I ), the setup described by Table 1 , and whose results are shown in Figure 4 , will represent the face-on view. To also consider the edge-on view, we perform another simulation where the initial magnetic field, instead of being homogeneous, is the CS given bywith B 0 = 3.0 G. The preparation of the fan is performed with the two stages as in the previous case, but in this case the shear (Equation ( 12 )) coincides with the CS location. The instantaneous pressure pulses are again applied at time t = 305.9 s. Figure 7 shows the temperature (i) and the density with overlapped magnetic field lines (ii) at an early time t = 312 s. The resulting pulse patterns are labeled ( E ), ( F ), ( G ), and ( H ). Figures 7 (c)–(d) show the same evolved patterns 1 minute later, at t = 372 s, for the temperature with magnetic field lines superimposed (iii) and the EM with temperature contours (iv). Same as before, the EM contrast of the SAD candidates will fade away fast and before the desired time if during their evolution they are not sufficiently isolated by the magnetic field. In fact, the feature ( F ) triggered in a location where the magnetic field lines are almost straight is not appreciated even at an early time (Figures 7 (a)–(b)). The cavities that survive and can be identified as SADs are located mainly in the neighborhood of the CS, 7 7 This is noted more clearly when the turbulence stage is carried out using larger wavenumbers compared with those chosen here. See also Paper I . where the magnetic field is rolled up, forming closed vortices and islands. Figure 7. (a) Temperature at t = 312 s; ( E ), ( F ), ( G ), and ( H ) labeled cavities triggered at t = 305.9 s. (b) Density at t = 312 s. Same as before, but with magnetic field lines superimposed. (c) Same as (i), but 1 minute later, at t = 372 s. (d) EM, also at t = 372 s. Cavities ( G ) and ( H ) have an EM
Paper: 2015MNRAS.452..878M (acknowledgments)
We thank Neal Katz for useful discussion and Phil Hopkins for comments on the manuscript. We thank the anonymous referee for a constructive report that helped improve the manuscript. CCH is grateful to the Gordon and Betty Moore Foundation for financial support and acknowledges the hospitality of the Aspen Center for Physics, which is supported by the National Science Foundation Grant No. PHY-1066293. PSB was supported by a Giacconi Fellowship provided through the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy under NASA contract NAS5-26555.
Paper: 2020ApJ...893...18A (fulltext)
H 0 has been proposed previously (Zibin et al. 2007 ). We take the additional step of propagating the uncertainties on T CMB to H 0 and present example noise curves to produce such a measurement, including a detailed treatment of foregrounds. Studies of the time evolution of the CMB angular power spectrum and forecasts for observations have also been conducted (Lange et al. 2007 ; Zibin et al. 2007 ; Moss et al. 2008 ). In particular, Lange Page ( 2007 ) found that the change in the power spectrum could in principle be detected over the course of a century. 2. Method 2.1. Hubble Cooling of <italic>T</italic><sub>CMB</sub> The monopole temperature of the CMB, hereafter denoted , decreases in time with the scale factor asThe time derivative of T is then . This can be evaluated to first order, taking as an example and T 0 = 2.725 K, 6 6 The forecast for depends on the true value of the current CMB temperature T 0 but the change in is small given the current sub-mK constraints on T 0 from COBE/FIRAS (Fixsen et al. 1996 , 2011 ). Therefore the temperature of the CMB decreases by 2 nK over a 10 yr period. One could thus in principle measure the real-time cooling of the CMB due to the expansion of space and from this infer the value of the Hubble constant. 2.2. Requirements to Measure <italic>H</italic><sub>0</sub> The simplest method to measure the Hubble constant from the monopole temperature would amount to fitting for the slope of T as it decreases in time. The model is linear in time with a slope of ,We have in mind a future version of the COBE/FIRAS experiment, where the data can be reduced in discrete time intervals to evaluate the monopole temperature of the CMB repeatedly over the mission duration. For example, the mission could take data continuously for 10 yr, potentially mapping out large portions of the sky, and the monopole CMB temperature would be inferred from these data every year. Given the expected level of cooling, we can then propagate uncertainties from T ( t ) to H 0 and determine what sensitivity to T ( t ) is required to measure H 0 with a given precision. We make some simplifying assumptions that will allow us to analytically calculate the uncertainty on H 0 using an ordinary least-squares estimator. First, we assume that T ( t ) is measured with (foreground-marginalized) uncertainty σ T in discrete uniform time intervals with spacing Δ t , such that t = k Δ t , where k = 1,2..., N . The mission duration is thus t total = N Δ t with N data points. Second, the uncertainties on T ( t ) are time-independent, parameter-independent, and Gaussian uncorrelated. 7 7 This is not entirely true as the CMB photon noise is one of the dominant noise sources and therefore the noise depends on T CMB . However, the change in the CMB photon noise over these timescales is negligible. The foreground-marginalized uncertainty σ T can then be written in terms of an effective sensitivity s T of the experiment such that,where s T is understood as
Paper: 2016ApJ...816..100J (acknowledgments)
The Fermi -LAT Collaboration acknowledges generous ongoing support from a number of agencies and institutes that have supported both the development and the operation of the LAT as well as scientific data analysis. These include the National Aeronautics and Space Administration and the Department of Energy in the United States, the Commissariat à l’Energie Atomique and the Centre National de la Recherche Scientifique/Institut National de Physique Nucléaire et de Physique des Particules in France, the Agenzia Spaziale Italiana and the Istituto Nazionale di Fisica Nucleare in Italy, the Ministry of Education, Culture, Sports, Science and Technology (MEXT), High Energy Accelerator Research Organization (KEK), and Japan Aerospace Exploration Agency (JAXA) in Japan, and the K. A. Wallenberg Foundation, the Swedish Research Council, and the Swedish National Space Board in Sweden. Additional support for science analysis during the operations phase is gratefully acknowledged from the Istituto Nazionale di Astrofisica in Italy and the Centre National d’Études Spatiales in France. Facility: Fermi -(LAT).
Paper: 2019ApJ...870..107D (fulltext)
has an outer radius of 38 kpc and Bruce Rout found a value of 55 kpc (Bruce 2009 ). Figure 1. Rotation curve of NGC 3198. Figure 1 shows the observed rotation curve of NGC 3198 with two fitting lines: the line “ab” for the radius r r t and the line “bc” for the radius r > r t . The quantities α (the slope of the line ab), β (the slope of the line bc), r t (the turn-off radius), and γ (the intersection point between the line bc and the velocity axis) were taken from the rotation curve and their values are shown below,The flatness characteristic of spiral galaxies is bound with the eccentricity ( k ) calculated from the relation (Binney Tremaine 1987 ). Calculations of the mass and density at different radii of the galaxy NGC 3198 using the above values taken from rotation curve ( α , β , γ , and r t ) and using Equations ( 8 )–( 11 ) were carried out using a computer program constructed for this purpose. The central inner density (at r r t ) calculated for NGC 3198 is found to be constant ( ρ 1 = 0.387 M o pc −3 ), where M o is one solar mass. The density of the outer regions ( r > r t ) of the spiral galaxy NGC 3198 is much less than that of the inner regions. The value of ρ 2 (r) changes from 0.0187 M o kpc −3 at r = 4 kpc to reach 0.00278 M o pc −3 at r = 30 kpc, so the outer density decreases with increasing the distance from the center of the galaxy (Figure 2 ). Figure 2. Outer density as a function of radius. The behavior of the calculated densities of the spiral galaxy NGC 3198 (constant inner density and declining outer density) is consistent with other works (Rubin et al. 1977 ; Binney Tremaine 1987 ; Daod et al. 2014 ). The inner mass (at r r t ) of spiral galaxy NGC 3198 grows up very quickly, as in Figure 3 . The inner mass grows from 1.334 × 10 9 M O at r = 0.1 kpc to a mass of 8.5391 × 10 9 M O at a distance of 4 kpc from the galactic center. Therefore, the mass at the turn-off radius (4 kpc) is more than 10 4 times the mass near the center of the galaxy. Figure 3. Inner mass distribution. The calculated outer mass of NGC 3198 changes from 8.539 × 10 9 M O at the turn-off radius to 5.262 × 10 10 M O at the outer radius ( r g = 30 kpc) of the galaxy. In other words, the mass at the edge of the galaxy is only five times more than that at the turn-off radius. Therefore, the outer mass ( r > r t ) of the galaxy NGC 3198 grows slowly compared to the inner mass, as in Figure 4 . Since outer borders of spiral galaxies extend to a large distance from the visible region. The outer halo mass is significant and it is four times more than the mass of the inner (more dense) regions, which confirm the existence of a massive halo and giving evidence as to the existence of dark matter (Begeman 1989 ). Figure 4. Mass distribution of NGC 3198. The galaxy is usually divided into four regions, the nucleus, the bulge, the disk, and the halo. The mass of these four regions can be calculated if the radius (distance from the center of the galaxy to the
Paper: 2019MNRAS.485.4841R (acknowledgments)
ACKNOWLEDGEMENTS This work has been supported financially by Research Institute for Astronomy Astrophysics of Maragha (RIAAM) under research project No. 1/6025-33.
Paper: 2020A&A...640A..74S (fulltext)
and B1 shock regions along the outflow of L1157 are marked here with circles (B0), triangles (B1), and squares (B2). The dashed circles represent the HPBW of the IRAM 30 m spectra at 86 GHz (29′′), 210 GHz (12′′), and 350 GHz(7.5′′). The green contours in the channel maps show the 10, 60, and 120 σ levels of the integrated intensity map of 28 SiO, shown in Fig. 1 . 4 LVG analysis In this section we analyse the single-dish spectra of the SiO and its isotopologues 29 SiO and 30 SiO from the 2−1 up to the 8−7 transition ( E up from 6 to 163 K), shown in Fig. 2 . The 12–11 spectrum was excluded from the analysis because it is too noisy: we note, however, that the fluxes reported in Table 1 are consistent with the solutions provided below. The line properties are summarised in Table 1 . This table also reports the integrated intensity in three velocity ranges: EHV: –23 to –13 km s −1 , HV: –13 to –3 km s −1 , and LV: –3 to +5 km s −1 . In order to derive the average physical properties of the gas that emits SiO, we carried out a non-LTE analysis using the large velocity gradient (LVG) code described in Ceccarelli et al. (2003) . We used the collisional coefficients of SiO with H 2 computed between 10 and 300 K for the first 20 levels by Dayou Balança (2006) that are provided by the BASECOL database ( Dubernet et al. 2013 ). We ran a large grid of models (≥7000) covering a SiO column density N SiO from 1.5 × 10 13 to 8 × 10 15 cm −2 , a H 2 density n H2 from 1 × 10 3 to 2 × 10 7 cm −3 , and a temperature T from 30 to 300 K. We then determined the solution by simultaneously fitting the 28 SiO, 29 SiO, and 30 SiO lines and left as free parameters N SiO , n H2 , and T , while the emitting sizes and line widths were those measured. We assumed 28 Si/ 29 Si = 19.7 and 28 Si/ 30 Si = 29.8, as obtained by Anders Grevesse (1989) and confirmed by our SOLIS observations for L1157-B1 (see Sect. 3.3 ), and we added the 20% calibration uncertainty in the observed intensity error bars. We sampled the parameter space unevenly (with steps of 10 K until 100 and 20 K for higher temperatures, steps of a factor of 2 for densities lower than 6 × 10 6 cm −3 and 1.3 at higher ones, steps of a factor of 1.5 in column densities until 1 × 10 14 cm −2 and 1.1–1.2 at higher ones to ensure better cover close to the minimum χ 2 . The beam dilution was taken into account, assuming that the emitting gas has a circular Gaussian shape and using the standard equation for beam dilution. We carried out the non-LTE analysis for the HV and EHV components. The results are described in the two following subsections and are summarised in Table 2 . Table 2 Result of the non-LTE LVG analysis. Notes. The HV and EHV emissions are composed of two physical components each: the first is associated with the B1a clump, and the second is more extended (see text). ( a ) SiO abundance with respect to molecularhydrogen. The H 2 column density used to calculate the SiO abundance is derived from the CO
Paper: 2015ApJ...799..209W (acknowledgments)
We thank the ESO staff, and in particular the staff at Paranal Observatory, for their helpful and enthusiastic support during the many observing runs over which the KMOS GTO were carried out. We thank the entire KMOS instrument and Commissioning team for their hard work, which allowed our observational program to be carried out so successfully. We also thank the software development team of SPARKfor all their work with us to get the most out of the data. This paper and the KMOS 3D survey have benefitted from many constructive, insightful, and enthusiastic discussions with many colleagues whom we are very grateful to, especially M. Franx, A. Renzini, K. Whitaker, K. Glazebrook, and R. Bassett. D.J.W. and M.F. acknowledge the support of the Deutsche Forschungsgemeinschaft via Project ID 387/1-1. We thank the referee for a thorough reading and valuable comments.
Paper: 2016ApJ...816....8H (acknowledgments)
The authors thank a referee, H. J. Melosh, for useful comments that substantially improve the quality of this work. Numerical computations were in part carried out on the GRAPE system at Center for Computational Astrophysics, National Astronomical Observatory of Japan. Y.H. is supported by an EACOA Fellowship that is supported by East Asia Core Observatories Association, which consists of the Academia Sinica Institute of Astronomy and Astrophysics, the National Astronomical Observatory of Japan, the National Astronomical Observatory of China, and the Korea Astronomy and Space Science Institute.
Paper: 2019MNRAS.485L..34B (fulltext)
The fall-off dependence of the dark-energy density as a function of r is given by the parameter n (c.f. eq. 1 in paper II). For a value of n = 3 the dark-energy density versus radius dependence does not contradict the tests of standard GR in the Solar system. Nielsen Birnholz ( 2018 ) have pointed out that a larger value of n is required to account for the results obtained from gravitational wave observations (c.f. their eq. 4). Paper II shows that for larger values of n , the predicted intensity and velocity structures will be almost identical within pc-GR, with only slight changes in the absolute values. For instance, increasing n from 3 to 4 results in a shift in the maximum orbital frequency from 1.72 to 1.89 m and a shift in the maximum orbital frequency from 0.21 to 0.22 (c.f. eq. 4 of paper II and Fig. 2 in this paper). Therefore the appearance of a dark and a bright ring in the intensity profiles discussed in Section 2 are robust and will not change for steeper values of n . In this paper we assume a value of n = 4 for the tests on the timing properties to account for the detection of gravitational waves. We assume a value of n = 3 for the ray-tracing simulations, as the structures in the intensity behaviour will not change. The Galactic Centre has been studied in great detail with near-infrared observations. Based on the orbital motions of stars in the vicinity of the Galactic Centre, the mass of the central black hole has been constrained to [Gillessen et al. ( 2009 ), see also Genzel et al. ( 1996 ) and Eckart Genzel ( 1997 )]. Gravity Collaboration et al. ( 2018a ) have, for the first time, measured the gravitational redshift of the star S2 near the Galactic Centre. In addition, Gravity Collaboration et al. ( 2018b ) have detected bright flares at distances of approximately 150 μas from the central black hole. In this paper we summarize our imaging and timing predictions based on the pc-GR theory for the Galactic Centre. In Section 2.1 the ray-tracing simulations are summarized. In Section 2.2 we predict the orbital velocities of particles which can be compared to Event Horizon Telescope (EHT) observations at different epochs. Finally in Section 2.3 we compare the emission profiles predicted by pc-GR and standard GR, which show remarkable differences. The theory of pc-GR in general, along with a specific focus on the tests discussed in this work, is presented in Hess et al. ( 2019 ) (hereafter Paper II : Theory and robust predictions). 2 OBSERVATIONAL TESTS In this section we describe the observational tests we have developed to date to verify the predictions of the pc-GR theory in preparation for the upcoming EHT observations. These tests include ray-tracing simulations of the flow of accreting matter around black holes using both standard GR and pc-GR. In addition, we have produced diagnostic plots comparing both theories with regard to emissivity profiles and the timing properties of orbiting matter. The main results are
Paper: 2017ApJ...845..151R (fulltext)
to resolving the Ba puzzle. To examine this possibility, we have selected a set of nine iron lines from the table in Gurtovenko Sheminova ( 2015 ) whose formation depth is far from the continuum forming layers in the range of log ( τ 5000 ) = −2.0 to −4.0. We show in Figure 12 the differential iron abundances of these lines (Table 4 ) covering the EWs range 60 to 160 mÅ in the spectra of HD 96116 and HD 222582 as a function average depth of line formation. All these abundances are measured using the stellar parameters listed in Table 1 whose values are based on the analysis of weak and moderate strength iron lines of EW 70 mÅ (Nissen 2015 ). It is evident from Figure 12 that HD 96116 with relatively high stellar activity has a systematically higher iron abundances for the strong lines formed in the upper layers of the photosphere (log ( τ 5000 ) −2.0) than the mean abundance of +0.01 dex derived for lines formed deep in the photosphere. Such a trend of iron abundance with depth of line formation strongly suggest a height-dependent microturbulence with velocity increasing toward the top of the photosphere and the transition to the chromosphere. From the very similar iron abundances for the strong Fe lines in HD 222582 suggest that the photospheric structure of the inactive star HD 222582 of age 7 Gyr is very similar to that of the Sun. Figure 12. Differential iron abundance derived from the adopted stellar parameters in Table 1 as a function of the average depth of line formation. Table 4 Iron Abundances for a Set of Nine Saturated Iron Lines along with the Average Depth of Formation in the Stellar Photosphere These results strongly suggest that the microturbulence increases steeply with height in the photospheres of active stars relative to dwarfs of age less than the Sun. As a result, the application of the microturbulence value derived from the weak and moderate strength iron lines to strong lines such as the Ba ii lines result in a serious overestimation of the Ba ii abundances in stars younger than the Sun. For the star HD 96116, an increase of microturbulence by about 0.3 km s −1 will drop [Ba/Fe] from the Ba ii 5853 Å line close to the abundance measured for the other heavy elements La−Sm. It is well understood from the analysis of the Sun that the chromospherically active phase of the Sun is associated with active upper photospheric layers (Schrijver et al. 1989 ) and thus chromospherically active young stars in the field, OCs, and associations have active photospheric layers from where the Ba ii lines emerge. As a result all those active stars, as found here and in previous studies, have Ba ii abundances overestimated by the standard LTE abundance analyses where the microturbulence is assumed depth-independent and derived from the weak/moderate strength Fe lines whose formation depth in the photosphere is very deep and differ greatly from the Ba ii line forming depths. Accordingly, we argue that the Ba abundances determined from the