adsabs/WIESP2022-NER huggingface dataset:
Paper: 2016MNRAS.463.4516P (fulltext)
because their terminal velocity increases due to the increasing gravitational field strength at small planetary radii. Only bolides in a very limited size range are large enough that they do not reach the sublimation temperature and yet small enough for their low terminal velocities to render them immune to frictional ablation. Given the uncertainties and limitations (for example, of excluding fragmentation and lateral spread) that enter our modelling, this narrow range may indeed be either somewhat smaller or simply non-existent in a scenario with more modes of ablation. If impacting planetesimals move with hypervelocities beyond the escape velocity of Jupiter of approximately 60 km s −1 , some high φ i bolides can ablate and still escape the upper outer envelope beyond a normalized altitude of 1. Fig. 6 shows the case of bolides impacting at right angles to the local zenith to 126 physical seconds. This time value is 10 times the time required for vaporisation of a 30 m bolide travelling at 100 km s −1 and is a good benchmark time to gauge when all highest velocity planetesimals either escape the system or completely ablate. Planetesimals capable of escaping the Jovian envelope are represented by black stars. Bolides of sizes 30–200 m with ∼80–100 km s −1 completely ablate upon impact, whilst larger bolides with these high kinetic energies manage to escape the Jovian radius (e.g. a 1 km planetesimal retains ∼0.60 of its initial mass in the process). At future times, the region of the plot involving velocities greater than the escape velocity evolves little whilst the low-velocity regime is subject to progressive ablation as in Fig. 5 . The former region does not evolve because we do not model the re-entry of planetesimals that may fall back to the planet due to their decreased velocities. Planetesimals of velocities exceeding Jovian escape offer chemical depositions that should be observable perhaps even in the atmosphere of such planets given the proximity of the 1 bar level and mixing processes. Figure 6. Mass ablation fraction to 126 s for planetesimals incident tangentially with initial velocities exceeding that of Jovian escape. Black stars represent escaped planetesimals. We see that planetesimals of velocities exceeding the Jovian escape should offer chemical depositions that should be observable at the 1 bar level. 5 DISCUSSION Our model is a simple description of planetesimal dynamics with implications for the enrichment of Jovian-like exoplanetary envelopes. There is consequently extensive potential for further improving the sophistication of our model and we here discuss a few examples of potential refinements. First, Nellis ( 1997 ) highlights that there is indeed a continual process of molecular dissociation into a metallic plasma; there is thus no real distinct boundary between molecular layers and monoatomic conductive hydrogen and the band-gap separation is bridged continuously with increasing pressure rather than jumping
Paper: 2020ApJ...904...51C (fulltext)
with Equation ( 5 ). From Figure 1 it is expected that data points at high enough accelerations do not have any sensitivity to e . Indeed, for data points with −10.3 x 0 −9 the distribution of Δ ⊥ gives a null result. Similarly, for galaxies with , the distribution of e (median value) gives a null result. Figure 4. Distributions of Δ ⊥ and e from fitting Equation ( 5 ) to the SPARC galaxies. The top panels show the distributions of orthogonal residuals Δ ⊥ for three acceleration bins from the MCMC results shown in the middle column of Figure 3 . The mean of the distribution is displaced from zero for lower acceleration bins, indicating declining RCs. The bottom panels show the distributions of the e values fitted to the individual galaxies binned by the median values of x 0 within the galaxies. As expected, for the galaxies in the high acceleration bin (), the data do not have any sensitivity to e and so the distribution has a mean of ∼0. For lower acceleration bins the distributions are shifted to positive e with high statistical significance, indicating a preference for the EFE. The broad distributions are due to the broad individual posteriors on e . Data points at low enough accelerations will have sensitivity to e and distributions with nonzero mean value are expected from Figure 3 . For data points with x 0 −11.3 the distribution of Δ ⊥ has a mean of −0.061 ± 0.008 (a bootstrap error) which is statistically significant at more than 7 σ . For a much larger number of data points with x 0 −10.3, Δ ⊥ has a smaller deviation of −0.035 ± 0.003, but the statistical significance of the deviation is more than 11 σ . Figure 5 shows individual e values and their uncertainties for a subset of 113 galaxies with median . Due to the large uncertainties on e , some galaxies can occasionally return negative values. However, the median value of e is 0.052 ± 0.011 (bootstrap error), which represents ≈5 σ detection of positive e . This value is statistically consistent with the median environmental gravitational field for these galaxies ( (bootstrap error)). Furthermore, based on the robust binomial statistic with equal probabilities for e > 0 and e 0, 78 cases of e > 0 out of 113 is 4 σ away from the expected mean of 56.5 cases. Figure 5. Statistical detection of the EFE. The median values of e from rotation-curve fits are compared with e env values from the large-scale galaxy environment (Desmond et al. 2018 ) for 113 galaxies with . The heat map considers the posterior probabilities of individual e measurements. The distributions of e and e env are shown by the top and right histograms. The median value of e is clearly offset from zero, indicating a 5 σ statistical detection of the EFE. The value of is statistically consistent with (see also Appendix B ). The individual galaxies considered in Figure 2 are indicated: for the golden galaxies at exceptionally high e env values, e is significantly different from zero at 8 σ (NGC 5033) and 11 σ (NGC
Paper: 2017MNRAS.472..483F (acknowledgments)
Acknowledgements The authors thank the anonymous referee for very helpful comments that improved this paper. The authors further thank A. Ferrara, R. Bouwens and N. Reddy for very helpful discussions related to this work. This paper makes use of the following ALMA data: ADS/JAO.ALMA#2013.1.00151.S, ADS/JAO.ALMA#2013.1.00034.S and ADS/JAO.ALMA#2012.1.00523.S. ALMA is a partnership of ESO (representing its member states), NSF (USA) and NINS (Japan), together with NRC (Canada) and NSC and ASIAA (Taiwan) and KASI (Republic of Korea), in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO and NAOJ. Based on data products from observations made with ESO Telescopes at the La Silla Paranal Observatory under ESO programme ID 179.A-2005 and on data products produced by TERAPIX and the Cambridge Astronomy Survey Unit on behalf of the UltraVISTA consortium. This work was supported by the Swiss National Science Foundation through the SNSF Professorship grant 157567 ‘Galaxy Build-up at Cosmic Dawn’. ES acknowledges funding from the European Research Council (ERC) under the European Union's Horizon 2020 research and innovation programme (grant agreement No. 694343). BG gratefully acknowledges the support of the Australian Research Council as the recipient of a Future Fellowship (FT140101202). MTS acknowledges support from a Royal Society Leverhulme Trust Senior Research Fellowship (LT150041). VS acknowledges support from the European Union's Seventh Frame-work programme under grant agreement 337595 (ERC Starting Grant, ‘CoSMass’). OLF acknowledges support from the European Research Council Advanced Grant ERC-2010-AdG-268107-EARLY.
Paper: 2019ApJ...870...32K (acknowledgments)
We thank Alberto Krone-Martins and Lennart Lindegren for expert advice on Gaia , Floor van Leeuwen for advice about astrometry, and Fred Adams, Michael Grudić, and Anushree Sengupta for useful feedback on the article. We would also like to thank the referee for helpful comments. E.D.F. and K.V.G. are supported by NASA grant NNX15AF42G and Chandra grants AR7-18002X, GO7-18001X, and G06-17015x, and the Chandra ACIS Team contract SV474018 (G. Garmire L. Townsley, principal investigators) issued by the Chandra X-ray Center under contract NAS8-03060. Facility: Gaia. Software: astro (Kelvin 2014 ), celestial (Robotham 2016 ), MVN (Korkmaz et al. 2014 ), R (R Core Team 2018 ), relaimpo (Grömping 2006 ), spatstat (Baddeley Turner 2005 ), TOPCAT (Taylor 2005 ).
Paper: 2019MNRAS.488.4740P (acknowledgments)
ACKNOWLEDGEMENTS We thank Indranil Banik and Tim Clifton for helpful discussions. CP has been supported by an STFC studentship.
Paper: 2020MNRAS.492.1257M (fulltext)
the LAEs in each band (AB); (4) lensing magnification factors derived in Shu et al. ( 2016c ) for the first four lensed LAEs, Marques-Chaves et al. ( 2017 ) for BG1429+1202, and in Section 3 for BG1501+3042; (5) stellar UV continuum slope (equation 2 ), UV dust extinction (AB mag, equation 3 ), and colour excess of the stellar continuum assuming the Reddy et al. ( 2015 ) attenuation curve; (6) observed flux (units of 10 −16 erg s −1 cm −2 ), rest-frame equivalent width (Å), and full width at half-maximum (km s −1 ) of the Lyα line; (7) centroid velocity of Lyα and ISM absorption lines (average) relative to the systemic redshift (km s −1 ); (8) low-ISM and H i covering fractions using the depth of the average absorption line profile of ISM lines (Si ii 1260 Å, C ii 1334 Å, and Si ii 1526 Å) and using the relation of Gazagnes et al. ( 2018 ), respectively; (9) intrinsic UV absolute magnitudes (AB) and star formation rates (M ⊙ yr −1 ; equation 4 ); (10) intrinsic Lyα luminosity (units of 10 42 erg s −1 ) and star formation rates (M ⊙ yr −1 ; equation 9 ); and (11) Lyα escape fraction. a The ACAM/WHT i -band data of BG0201+3228 are affected by a high background level. The projected separation between the centroids of the two components in the source plane is about 0.46 ± 0.06 arcsec which at z = 2.645 corresponds to 3.8 ± 0.6 kpc for the adopted cosmology, suggestive of a close interacting/merging system. Nevertheless, we cannot confirm unambiguously the merging scenario as no velocity offset between the lensed images ‘A’ and ‘B’ is clearly detected in the available spectrum (with a resolution of ≃ 450 km s −1 FWHM). The sources have effective radius kpc and kpc. 4 RESULTS FROM IMAGING DATA 4.1 Photometry As seen in Fig. 1 the morphology of the LAEs in this sample consists of complex arc-like features with multiple lensed images seen in the observer plane. Given the typical ∼0.8 arcsec FWHM seeing of the GTC, WHT, and CFHT images, there is light contamination from the foreground lens galaxies. In order to perform a clean photometry of the lensed images, we use the two-dimensional fitting program galfit (Peng et al. 2002 , 2010 ) to model and substract the light distribution of the foreground lens (and in some cases additionally nearby objects). We use Sérsic profiles centred at the centroids of each LRG light emission allowing only two pixels freedom (∼0.5 arcsec). For the four lensed LAEs with HST imaging, we use the models of the foreground lens presented in Shu et al. ( 2016c ) as priors in our modelling (i.e. effective radius, minor-to-major axial ratio, and major-axis position angle). Nearby stars were used as PSF models. We further use SExtractor (Bertin Arnouts 1996 ) to measure the magnitudes in each band. To do so, we use aperture photometry using apertures with diameters of 2.5 times the FWHM of the PSF. Aperture photometry was performed for compact lensed morphologies, such as the lensed images of BG0755+3445, BG1429+1202 and the
Paper: 2017ApJ...837..175S (fulltext)
line), along with the combined contribution (black dashed line) as given in van Soelen et al. ( 2012 ). A 33,000 K blackbody distribution is plotted for comparison (blue dashed line). 3. Gamma–Gamma Absorption Gamma-ray photons emitted in binary systems are subject to γγ absorption as they pass through the star’s photon field (see, e.g., Dubus 2006 ). The interaction of a γ -ray photon with a low-energy photon can result in electron–positron pair production (Gould Schréder 1967 ), if the energy exceeds the threshold condition for pair production, which is given bywhere and denote the photon energies of the target low-energy photon and the γ -ray photon, respectively, normalized to the electron rest-mass energy, i.e., , and θ is the interaction angle between the two photons. The minimum threshold occurs for a head-on collision (), and for a 1 TeV gamma-ray photon this will require photons with a threshold frequency ofwhich is within the mid-IR regime (where the disk is the primary photon source; Figure 2 ). The γγ optical depth is given by (Gould Schréder 1967 )where l is the distance over which the γ -ray photon travels, , , and is the number density of the low-energy target photons. Here, is the γγ cross section (Jauch Rohrlich 1976 ),whereand is the Thomson cross section. The maximum γγ cross section occurs when the energy of the interacting photons is twice that of the energy threshold (Equation ( 8 )). This implies that for TeV γ -ray photons the γγ interaction is a maximum for IR photons at a frequency We treat the two components of the target radiation field (stellar radiation and circumstellar disk radiation) separately and numerically calculate the γγ optical depth of both components for a γ -ray photon traveling in the direction of the observer. The total optical depth is given by the sum of the optical depths of the two components. Figure 3 shows the optical depth for a γ -ray photon with an energy of 1 TeV as a function of time from periastron. As expected, the γγ optical depth due to the photons from the disk is lower than that due to the star but still results in a decrease of the flux at TeV energies. The total decrease of the flux at 1 TeV due to γγ absorption may reach ∼60%. For the geometrical configuration used in our calculations (see Section 2.1 ) the maximum absorption due to the disk is at ≈4 days prior to periastron and that due to the star is at ≈2 days before periastron. Figure 3. The γγ optical depth of the Be star’s radiation field as a function of time around periastron for a 1 TeV γ -ray photon. The blue dashed line denotes the stellar component, the red dot-dashed line denotes the circumstellar disk component, and the combined effect is shown by the solid black line. 4. TeV Emission We assume that the TeV emission from PSR B1259-63/LS 2883 is produced by relativistic particles accelerated at the termination shock that occurs between the pulsar and stellar winds. The location of the termination shock depends on the
Paper: 2016MNRAS.460...19M (acknowledgments)
This research has made use of data obtained from the SIMBAD data base operated at CDS, Strasbourg, France; the High Energy Astrophysics Science Archive Research Center (HEASARC), provided by NASA's Goddard Space Flight Center NASA/IPAC Extragalactic Database (NED). The authors acknowledge the ASI financial support via ASI–INAF grant 2013-025-R.0.
Paper: 2016ApJ...833....2G (fulltext)
novel prediction for the scatter in the mass in the stellar halos. Upcoming surveys (e.g., van Dokkum et al. 2014 ) of stellar halos will help to constrain this picture. The scatter of the SMHM relation contains valuable information on the galaxy formation process. It deserves careful consideration in both observations and simulations in the future. At the high-mass end (), the scatter is a generic consequence of the hierarchical assembly. Different simulations should be able to produce consistent scatter of the SMHM relation in this mass range. At lower halo mass () the scatter gradually becomes dominated by baryonic processes involved in the in situ buildup, and has potentially more constraining power on the underlying theory. Two different processes, in situ and ex situ growth, are needed to produce the apparent flat scatter in the SMHM relation. However, the similar magnitude of the observed scatter at the low and high-mass end is likely a coincidence. Other works (e.g., Boylan-Kolchin et al. 2012 ; Garrison-Kimmel et al. 2014 ; Dutton et al. 2015 ; Sawala et al. 2015 ) have shown that at the regime of dwarf galaxies, the scatter of the SMHM relation may be much larger than that in more massive systems. More stringent observational constraints on the mass dependence of the scatter would be valuable. We end this section by highlighting several caveats and areas for future improvement. The Bolshoi simulation has a limited resolution (one particle mass ), so halos with mass smaller than may not be properly described in the Bolshoi simulation. As shown in Figure 1 , a fraction of satellite galaxies are possibly affected by this limited resolution. The result for host halos at the low-mass end () should be treated with caution. However, as discussed earlier, since the stellar component in such halos should be very small, based on the SMHM relation, this should only have a small effect on our result. Moreover, we adopt a simple assumption that when a subhalo disrupts, all of its stellar content is accreted to the host halo, while in reality, the accretion of stellar mass could be delayed (Wetzel White 2010 ). We also do not distinguish between the stellar mass in the central galaxy and that in the ICL. Since most observations only include the stellar mass measurement of the central part of the host halo, a direct comparison with observations should also be treated with caution. The scatter in the SMHM relations observed should be roughly the quadratic sum of the intrinsic scatter discussed above and the measurement error. Behroozi et al. ( 2010 ) carefully studied different uncertainties in the SMHM relation. The authors included statistical uncertainties in the stellar mass function, cosmological parameters, and the uncertainty within the methodology to construct the SMHM relation, and revealed a ∼0.07 dex (and ∼0.12 dex) observational uncertainty at z = 0 (and z = 1). As shown in Figure 8 , we compare this with the observed scatter at z = 0
Paper: 2019A&A...624A..78D (acknowledgments)
Acknowledgements E.D.M., V.Zh.A., M.T., N.C.S., and S.G.S. acknowledge the support from Fundação para a Ciência e a Tecnologia (FCT) through national funds and from FEDER through COMPETE2020 by the following grants UID/FIS/04434/2013 POCI–01–0145-FEDER–007672, PTDC/FIS-AST/7073/2014 POCI–01–0145-FEDER–016880 and PTDC/FIS-AST/1526/2014 POCI–01–0145-FEDER–016886. E.D.M. acknowledges the support by the fellowship SFRH/BPD/76606/2011 funded by FCT (Portugal) and by the Investigador FCT contract IF/00849/2015/CP1273/CT0003 and in the form of an exploratory project with the same reference. V.Zh.A., N.C.S. and S.G.S. also acknowledge the support from FCT through Investigador FCT contracts IF/00650/2015/CP1273/CT0001, IF/00169/2012/CP0150/CT0002 and IF/00028/2014/CP1215/CT0002 funded by FCT (Portugal) and POPH/FSE (EC). A.M. acknowledges funding from the European Union’s Horizon 2020 research and innovation program under the Marie Sklodowska–Curie grant agreement No 749962 (project THOT). W.J.C. acknowledges support from the UK Science and Technology Facilities Council (STFC). Funding for the Stellar Astrophysics Centre is provided by The Danish National Research Foundation (Grant agreement no.: DNRF106). J.I.G.H. acknowledges financial support from the Spanish Ministry project MINECO AYA2017-86389-P, and from the Spanish MINECO under the 2013 Ramón y Cajal program MINECO RYC-2013-14875. A.C.S.F. is supported by grant 234989/2014-9 from CNPq (Brazil). Finally, we thank the anonymous referee, whose detailed review helped to improve this paper. This work has made use of data from the European Space Agency (ESA) mission Gaia ( https://www.cosmos.esa.int/gaia ), processed by the Gaia Data Processing and Analysis Consortium (DPAC, https://www.cosmos.esa.int/web/gaia/dpac/consortium ). Funding for the DPAC has been provided by national institutions, in particular the institutions participating in the Gaia Multilateral Agreement. This research has made use of the SIMBAD database operated at CDS, Strasbourg (France). The core analysis was performed with R version 3.3.1 ( R Core Team 2016 ), RStudio Version 1.0.143, and the R libraries dplyr 0.5.0 ( Wickham Francois 2016 ), foreach 1.4.3 ( Analytics Weston 2015 ), corrplot 0.86 ( Wei Simko 2017 ), parallel ( R Core Team 2016 ), RcppParallel ( Allaire et al. 2016 ), and relaimpo 2.2-2 ( Grömping 2006 ).
Paper: 2018MNRAS.475.2642K (acknowledgments)
ACKNOWLEDGEMENTS We thank Åke Nordlund for his contributions to the development of the zoom-technique and the modified ramses version used at the Centre for Star and Planet Formation in Copenhagen, and for valuable suggestions and comments to the article. We thank the developers of the python-based analysing tool yt http://yt-project.org/ (Turk et al. 2011 ). Their efforts in providing analysis support for ramses simplified and improved our analysis significantly. We are grateful to the anonymous referee for constructive feedback that helped to improve the manuscript. This research was supported by a Sapere Aude Starting Grant from the Danish Council for Independent Research to TH. Research at Centre for Star and Planet Formation is funded by the Danish National Research Foundation (DNRF97). We acknowledge PRACE for awarding us access to the computing resource CURIE based in France at CEA for carrying out part of the simulations. Archival storage and computing nodes at the University of Copenhagen HPC centre, funded with a research grant (VKR023406) from Villum Fonden, were used for carrying out part of the simulations and the post-processing.