adsabs/WIESP2022-NER huggingface dataset:
Paper: 2019MNRAS.485.2505T (fulltext)
(Planck Collaboration I 2016 ), is implemented in order to constrain the Λ( H ) models. Up to now, for these models only the set of the CMB shift parameters have been used in order to test the dynamical vacuum scenarios close to recombination. Then, we combine CMB in an overall likelihood analysis with other cosmological probes (SNe type Ia, BAOs, H 0 ) in order to extract the probability distribution (via Monte Carlo MCMC method) of the solutions for a large set of cosmological parameters, as well as, to search for deviations from the concordance ΛCDM model for which Λ( H ) = const. The structure of our paper is as follows: in Section 2 we provide a brief introduction of the running vacuum cosmology and we present the most popular Λ( H ) models that have appeared in the literature. In Section 3 , we discuss the methodology and the cosmological data that we utilize and we perform a detailed statistical analysis aiming to provide the corresponding best-fitting values and contour plots for the current Λ( H ) models. Finally, we discuss our results in the conclusions. 2 BACKGROUND EXPANSION IN RUNNING VACUUM MODELS In this section, we briefly describe the main features of the dynamical vacuum models for which Λ is not constant but evolves with cosmic time. This is perfectly allowed by the cosmological principle embedded in the FLRW metric (Perico et al. 2013 ). In general, if we model the expanding Universe as a mixture of perfect fluids N = 1, 2, ... then the total energy momentum tensor is given by where is the four-velocity field. In this case the components of are written as where the quantities p T and ρ T are the total pressure and energy density in the comoving frame , respectively. The next step is to apply the covariant local conservation law for the mixture, namely ∇ μ T μν = 0. Inserting equation ( 1 ) into the latter expression and with the aid of the relation (which comes from the fact that for any four-velocity vector, we have ) we arrive at (Grande, Pelinson Sola 2009 ): For a Friedmann–Lemaître–Robertson–Walker (FLRW) metric, it is easy to check that for a comoving frame (), one obtains: and thus equation ( 3 ) reduces to where is the Hubble parameter and a ( t ) is the scale factor of the Universe normalized to unity at the present epoch. For the rest of the paper we focus on the spatially flat FLRW metric. In the aforementioned discussion, we did not address the physics of the fluids involved. The total density ρ T receives contributions from non-relativistic matter (cold dark matter and baryons) ρ m = ρ b + ρ dm ( p m = 0), radiation ρ r ( p r = ρ r /3) and vacuum ρ Λ ( p Λ = −ρ Λ ). In this context, the index N specifies the specific components of the cosmic fluid, namely {dm, b, r , Λ}. Assuming that baryons and radiation are self-conserved, namely the corresponding densities evolve in the nominal way ρ b = ρ b0 a −3 and ρ r = ρ r0 a −4 the overall conservation law ( 5 ) becomes or Note that Q is the interaction term
Paper: 2018MNRAS.473.3394V (fulltext)
have discussed the mechanisms that lead to the formation of instabilities in both the upstream and downstream medium when the magnetic field makes a large angle with the direction of the flow. We should note here that when we increase the angle even further beyond 70° this effect disappears since the component of the particle velocity parallel to the flow becomes smaller than the drift velocity if the thermal plasma, which eliminates the upstream current. The question remains as to why this result was not obtained with a hybrid code. By comparing our simulations with these earlier models we come to the conclusion that there are two factors that contribute to the different outcome. The number of particles . Although the absolute number of particles used in hybrid codes is far larger than what we use for the PI[MHD]C simulations, only a small fraction of the particles in hybrid simulations actually are suprathermal. Combined with the fact that for an oblique magnetic field only a fraction of the non-thermal particles are able to move upstream because of their pitch angle, we conclude that the hybrid simulations lacked the necessary number of particles to create the current that causes the disturbance of the upstream magnetic field. The size of the physical domain . The size of the simulated domain in the direction perpendicular to the shock-normal is six times larger in our model than in the original hybrid simulation. This allows us to capture the long wavelength that corrugates the shock and sets in motion the process that disturbs the upstream magnetic field. The smaller domain used by the hybrid simulation makes this impossible. We have tested our code with a physical domain similar to that used by Caprioli Spitkovsky ( 2014a ) and find that indeed the long wavelength never appears and so is the magnetic field amplification. These points demonstrate the opportunity that the PI[MHD]C method affords us. Because it is designed primarily to simulate large-scale structures, is less computationally expensive than PIC or hybrid methods and enables us to run the simulation in the frame of reference of the shock, it allows us to simulate a large volume of space, using a larger number of suprathermal particles, which in this particular case proved to be necessary. 2D3V PIC simulations of perpendicular and oblique shock configurations have been performed by Kato Takabe ( 2010 ) and Wieland et al. ( 2016 ), respectively. These studies verify that fast Alfvénic Mach (low magnetization) number shocks are mediated by the Weibel instability. These studies only report on electron thermalization but find pieces of evidence for the development of a proton non-thermal tail possibly connected with a shock surfing process in the perpendicular shock configuration. We confirm with Caprioli Spitkovsky ( 2014a ) that protons are accelerated at oblique shocks with an angle θ B = 45°. However, we argue that particle acceleration also proceeds at larger angles through
Paper: 2021MNRAS.501.4903F (fulltext)
of the local inventory, nearby solar-type stars provide particular insight from high signal-to-noise, high-resolution spectroscopy, such as stellar masses and ages. Most or possibly all of the scatter in the barium-to-iron abundances among the Population II and intermediate-disc stars, as we learn in this work, is subject to former mass transfer from white dwarf companions, and, if accounted for, discloses a very different perspective on the underlying core-collapse supernovae nucleosynthesis. Likewise, non-single sources cause false stellar parameters and hence scatter or potentially even false trends in the derived elemental abundances, if their individual components are not treated as such, which is often commonplace. There are numerous examples in the literature, and sources like the intermediate-disc star HD 117635 in P1 or the Population II star HD 34297 in Fig. 1 below in this work, just two of them. On the other hand, it is frequently found that towards the metal-rich end the Population I and II disc star sequences progressively coalesce with their iron versus α-chain elemental abundances. This is – for many reasons – important to know, and it may be of some relevance that the stars of P1 and the present, extended sample, as we see, do not comply with a systematic trend of this kind in terms of their magnesium and iron enrichment. Figure 1. Modelling of the Fe i line λ6200.321 in the composite spectrum of the spectroscopic binary HD 34297 ( P ≃ 2 yr) with its Aa and Ab components being Doppler-displaced at Δ v r (Aa–Ab) =+28.9 km s −1 in this 2017 November 28 observation. Stellar parameters for the spectrum synthesis (dark blue) are as given in the legend. The light blue-dotted curves denote abundance changes of Δ[Fe/H] = ±0.10 dex. Individual line profiles of the Aa and Ab components are given by the red-dotted curves. In Sections 2 and 3 we briefly discuss the characteristics of the underlying spectroscopic observations and the model atmosphere analyses. The very centrepiece of our work follows in Section 4 with the analyses and related key aspects of the Population II and intermediate-disc stars. As in P1, these comprise about one-third of the final 42 pc ancient star sample and several findings as they become apparent at the present, intermediate stage of the survey will be the subject of Section 5 . Population I stars, which in the course of the analyses are identified as such, and hence not part of the survey, are presented in the Appendix. 2 OBSERVATIONS With reference to a Gaia -based 42 pc parent sample, our strategy to cover an unbiased, complete subset of more than a hundred ancient solar-type Population II (τ ≥ 12 Gyr) and intermediate-disc stars (τ ≃ 10 Gyr) within that volume has been described in detail in P1. The spectroscopic observations that we present in this second work on the survey were all secured in 2017 with the FEROS échelle spectrograph (Kaufer et al. 1999 ) at the La Silla Observatory in Chile, most of them
Paper: 2016ApJ...817...88Z (fulltext)
be plausible for an isolated black hole accreting from the interstellar medium (see, e.g., Maccarone ( 2005 ) and references therein). At the same time, in this context there is an indication that the observed X-ray spectrum of source XMM0838+24 does not coincide with predictions for various forms of radiatively inefficient accretion flows that are expected to have a power law spectrum with a photon index 1.4–2.5 (see, e.g., Narayan et al. 1998 ; Ball et al. 2001 ). High-resolution X-ray and deeper optical/NIR imaging and spectroscopic observations are required to shed light on the association of XMM0838+24 and its host galaxy and thus derive its X-ray luminosity. In this study we do not attempt to analyze X-ray spectra of XMM0838+24 and XMM1226+12 in detail because given their faint X-ray fluxes it would not provide any selectivity to discriminate conventional X-ray object types. However, we compare their X-ray colors from the XMM-Newton source catalog to a sample of known AGNs and ESO 243–49 HLX–1. In Figure 6 we overplotted two “corner” states of ESO 243–49 HLX–1 and our HLX candidates on hardness diagrams of the known AGN population from Lin et al. ( 2012 ). Clearly, broadband properties of ESO 243–49 HLX–1 do not allow its unambigous separation from AGNs. At the same time, selected HLX candidates are situated aside from ESO 243–49 HLX–1 and most AGNs due to systematically harder HR2 and HR3 colors. This, however, does not favor any particular interpretation of the nature of the candidates. Figure 6. The X-ray color–color diagrams for known AGN (from Lin et al. ( 2012 ), black dots), selected HLX candidates (this paper, gray stars) and two spectral states of ESO 243–49 HLX–1, high luminosity consistent with steep power law and low luminosity hard power law states (gray triangles). X-ray colors are taken from the 3XMM-DR5 catalog and are defined as HR , where f i is the X-ray flux in i -th XMM-Newton band (band 1 is 0.2–0.5 keV, band 2 is 0.5–1.0 keV, band 3 is 1.0–2.0 keV, band 4 is 2.0–4.5 keV, band 5 is 4.5–12 keV). As claimed by Lin et al. ( 2012 ) systematic uncertainties of these colors can reach 0.1 which is below symbol size for ESO 243–49 HLX–1 and HLX candidates. 5.2. Base Sample Properties We compute the number density of the HLX candidates from our base sample. From a total of 98 HLX candidates we expect to have at least 21 true HLXs. Given the searched volume and solid angle this gives us an HLX number density of 1 object per ≈2 × 10 5 Mpc 3 . This is approximately a factor of several thousand less frequent than ULXs (Swartz et al. 2011 ). We can normalize this density to galactic halos. In total there are 33,879 SDSS galaxies falling in the 3XMM-DR5 footprint (total observed part of the sky) that we included in our off-nuclear cross-match procedure. Hence, the density of our HLX candidates is about 6 × 10 −4 per galaxy. If they are all accreting IMBHs, this is too large a density for a population of “naked” IMBHs that
Paper: 2017MNRAS.470.2240F (fulltext)
drive heat conduction in the disc. Fig. 6 shows that, even in 3D, the heat flux is negligible in regions in which P saturates. From our EMHD simulation, we thus deduce that the average impact of the non-ideal effects in the compact, SANE disc studied here might be close to what one would obtain by setting P 0.25 b 2 and q 0. We also note that our results critically depend on the treatment of the mirror instability in the EMHD model. We assume that when driven above the threshold of the mirror instability, P saturates at nearly exactly the instability threshold P mirror , through an increase in the effective collision rate in the plasma. Our results are only physical if this approximation captures the global behaviour of the plasma on length-scales much larger than the Larmor radius. While current particle-in-cell simulations broadly support this model (Kunz etal. 2016 ), comparisons of the EMHD model with more detailed local simulations of a plasma at the mirror instability threshold will be required to test this assumption. 4.4 Stress tensor and angular momentum transport In previous sections, we showed that the pressure anisotropy in the EMHD simulations of SANE discs is, in a time-averaged sense, P b 2 /4. In Section 2 , we also argued that the associated shear stress, , is effectively of the same form as the magnetic shear stress, . As far as the shear stress is concerned, the pressure anisotropy P and the magnetic energy density b 2 are interchangeable. Additionally, for 1, the magnetic energy density and isotropic magnetic pressure do not play a particularly important role in the dynamics of the accretion flow. Accordingly, we may expect similar steady-state accretion flows to be reached in EMHD and IMHD if . Our simulations appear consistent with this interpretation of the impact of the pressure anisotropy: in the SANE configuration, we find similar density and temperature profiles for the EMHD and IMHD models, but with a weaker magnetic field in the EMHD model, e.g. the plasma parameter satisfies EMHD (1.21.4) IMHD . We can have a more direct look at the component of the stress-energy tensor responsible for angular momentum transport in the disc, . We show radial profiles of (averaged over time and azimuthal angle, integrated over the polar angle and normalized by the gas pressure) in Fig. 10 . The first notable result is that the total stress is nearly identical for the IMHD and EMHD simulations. However, in the IMHD model, the entirety of the stress is due to the magnetic field. In the EMHD model, outside the ISCO of the black hole, (7075)percent of the r component of the stress is due to the magnetic field and (2530)percent to the pressure anisotropy. While this fraction of the stress due to the pressure anisotropy is slightly larger than one might have expected based on the simplistic model of a constant ratio P b 2 1 4, the properties of the stress tensor are broadly consistent with the view that in EMHD the pressure anisotropy
Paper: 2020MNRAS.491.5317I (fulltext)
have this feature, because we masked them out during the second iteration of the PSF construction process. As our PSF model cannot model all the particular features of a given individual star, after the subtraction process, some residuals are left. They correspond mainly to the central region due to the saturated pixels and the internal reflections. In the same way, the diffraction spikes of each star will remain as residuals since the PSF models do not have diffraction spikes. All of these effects can be seen in Fig. 10 , where we show a colour composite ( g, r , and i ) image of the central part of the Coma Cluster field in three different panels. Top panel shows the original image with the scattered light of near and bright stars contaminating the central parts of the galaxy cluster. In the middle panel, we show the scattered light field model of the brightest stars. In this particular case, we have modelled all stars brighter than 13 mag in the G band of Gaia DR2 Catalogue. The bottom panel shows the stars subtracted image. After the removal of the scattered light field, the low-surface brightness features of the Coma Cluster and its intracluster light are ready to be explored. Figure 10. Colour composite image of the central region of the Coma Cluster as seen by SDSS (R.A. (J2000) = 12:59:48.658, Dec (J2000) = +27:58:49.443), using g, r , and i filters. The background of the image (in grey) corresponds to the sum of g, r , and i SDSS band images. The images have been rebinned to 2 arcsec per pixel and convoluted with a Gaussian kernel of 3 pixel width to enhance the low-surface brightness features of the data. Upper panel: SDSS original image. Note the enormous contamination produced by the brightest stars in the central part of the Coma Cluster. Middle panel: scattered light field of all stars brighter than 13 mag in the G band of Gaia DR2 Catalogue. Bottom panel: scattered light field subtracted image. 4 DISCUSSION AND CONCLUSIONS We present extended SDSS PSFs (8 arcmin in radius) for all five SDSS filters. We describe the methodology we have used to obtain these large PSFs and we show a particular example of how the PSFs can be used to remove the scattered light field produced by the brightest stars in the images. Our approach of deriving PSFs can be applied to other telescopes and instruments to obtain extended PSFs. This work shows that the PSFs have complex bidimensional structures as it can be seen in Fig. 6 . In that figure we plot the radial profiles produced by de Jong ( 2008 ) with a maximum radius of ∼46 arcsec for the g, r , and i SDSS filters together with the radial profiles of the PSFs of this work. The reader can see that the agreement between these two works is rather good. Compared to de Jong ( 2008 ), we enlarge the model of SDSS PSFs an order of magnitude in radius. In addition to this, we also provide their 2-dimensional structures. The PSFs models presented in this work can be used to remove the scattered light of
Paper: 2020A&A...642A..86T (fulltext)
than the canonical value of 1.4. Fig. 12. Parallax over error ( π / σ π ) ve. RUWE of the three new groups of YSOs identified in this work, compared to sources with resolved protoplanetary disks (DSHARP, Andrews et al. 2018 ) and known binaries in Upper Sco ( Tokovinin Briceño 2020 ). 5. On the origin of the new YSOs with disks The previous section has shown that the new disk-bearing YSOs that are candidate members of the US and V1062 Sco comoving groups are relatively well characterized with respect to proper motion, on-sky spatial distribution, distance distribution, and age. Tables D.1 and E.1 list the source designations and disk type for US and V1062 Sco, respectively. The remaining YSOs in our sample (Table A.1 ) are likely composed of sources from both Lupus and UCL, as they present properties that are common to both of these populations; the emerging picture is that the youngest of these YSOs are likely Lupus members, while the older sources are likely part of UCL. Comparison of our new disk-bearing sources in Table A.1 with the YSOs identified by Damiani et al. (2019) yielded 37 out of 60 matches: 45% of the sources were found to belong to the UCL diffuse group D1, and 35% to the UCL diffuse group D2 (further divided into 27% from the D2a sub-group and 8% from the D2b sub-group). Of our new disks in US, 12/19 were identified as belonging to either the diffuse population in Upper Sco USC-D2 (32%), to UCL-D1 (32%), or to UCL-D2a (5%). Regarding the new disks in V1062 Sco, 14/19 were also previously identified by Damiani et al. (2019) as candidate members of the V1062 Sco/UCL-1 compact group. We discuss in this section the region morphology and kinematics to explore the origin and connection between Lupus, UCL, US, and V1062 Sco. 5.1. Evolution of the 3D spatial distribution We attempt to reconstruct the 3D distribution of the new disk-bearing YSOs and the Lupus clouds to better determine how the Lupus and UCL stellar groups are related. Figure 13 shows in the upper (middle) row the Galactic Latitude (Longitude) as a function of distance for these sources. Each row has three panels, corresponding to the three iscohronal age groups. These panels also include a schematic representation of the Lupus clouds (the central Galactic coordinates were taken from Hara et al. (1999) and the width of the filled ellipses are merely representative and not to be taken to scale). The distances to Lupus 1, Lupus 2, Lupus 3, and Lupus 4 were derived from the distances to their previously known members ( Dzib et al. 2018 ), and for clouds with no known associated YSOs we use the distances calculated by Zucker et al. (2019) (we exclude the clouds Lupus 7 and Lupus 8, due to the lack of distance estimates). To constrain the origin of a YSO, we explore its travel radius, specifically, the projected distance traversed by a source since its birth assuming a 1D velocity of 1 km s −1 . We represent the area corresponding to the travel radius as dotted-line
Paper: 2019ApJ...871..215B (fulltext)
is a potential avenue for direct comparison of the observations to the predictions of MHD simulations. The magnetic field lower limits found in each region are consistent with the values found by H13 . This peak value is consistent with the H13 value of +8 μ G. Our data strengthen their finding that the RMs along the narrow, decelerated H α ridge in region 5 are much stronger and opposite in sign compared to the surrounding RMs, suggesting a compressed ambient field which has a preferred direction. Due to the high density of RMs on and next to the Smith Cloud, we can model the morphology of the magnetic field. As seen in Figures 5 and 6 , an arch of positive RMs extends above and around the head of the Smith Cloud. This cap of RMs suggests a magnetic field that has been compressed by the ambient field as the Smith Cloud travels toward the Galactic plane, assuming that the cloud is falling into the disk (see Section 1 ). Under the assumption that is roughly constant, since RM is proportional to , the small negative RMs over the body suggest that the perpendicular component of the magnetic field is stronger than the parallel component, while the large positive RMs surrounding the cloud indicate that the parallel component dominates. To the right of the major axis of the cloud, the RMs are predominately positive, indicating a weak magnetic field perpendicular vector component pointing toward the observer. To the left of the Smith Cloud, strong negative RMs dominate. These strongly negative RMs indicate weak magnetic field perpendicular vector components pointing away from us. This picture is qualitatively consistent with strong parallel magnetic field vector lines along the sides of the Smith Cloud. This suggests a magnetic field that is draped, or laid over the ionized gas of the Smith Cloud, where the field has been compressed near the tip of the cloud (Dursi Pfrommer 2008 ). This configuration of a draped magnetic field surrounding the Smith Cloud is consistent with the idea that the neutral gas in the Smith Cloud is shielded from the hot ISM by a magnetic barrier between the surrounding medium and the gas; this has been seen in 2D simulations by Konz et al. ( 2002 ) and in 3D simulations by Grønnow et al. ( 2017 ). Grønnow et al. ( 2017 ) found that a magnetic field of a few μ G associated with an HVC within 10 kpc of the disk corresponds to an ambient field that has been “swept up” by the cloud, which is consistent with H13 and our observations. The Grønnow et al. ( 2017 , 2018 ) models also support the idea that a magnetic field could keep the cloud intact as it travels through the Galactic halo by delaying hydrodynamic instabilities and stripping of the gas. With this draped magnetic field barrier, the Smith Cloud could travel through the ISM while keeping the gas near the head and body intact. To test if this geometry could produce the observed RMs, we create a toy model of a possible magnetic field configuration surrounding the