adsabs/WIESP2022-NER huggingface dataset:
Paper: 2016A&A...586A.133P (fulltext)
Collaboration XI (2014) . This map is corrected for zodiacal emission and the brightest extragalactic point sources are subtracted. The Galactic reference offsets of the underlying IRAS and Planck data were obtained through a method based on correlation with 21 cm data from the LAB H i survey ( Kalberla et al. 2005 ) integrated in velocity, effectively removing the CIB monopole in the model map. The mean dust intensity, , listed for each LR region in Table 1 , ranges from to for increasing . The mean column density calculated from the LAB survey data is also listed in Table 1 , with values ranging from to . For all of the LR regions, we fit the and spectra with a power law in , using the fixed exponent , over the -ranges defined in Sect. 4.2 . The amplitudes derived from these fits are listed in Table 1 ; the amplitudes can be retrieved from the ratio. These amplitudes are plotted as a function of in Fig. 4 after normalization by the maximum value found for the largest region (LR72). We fit the empirical dependence of these amplitudes on as a power law of the form where . The two fitted exponents are quite similar, and . The exponent that we find for polarization is close to the one observed in the diffuse interstellar medium for the dust intensity, consistent with , where is the mean value of the dust specific intensity ( Miville-Deschênes et al. 2007 ). Values close to 2 are expected, because we compute angular power spectra, which deal with squared quantities. Although the data points roughly follow this dependence, the empirical law fails to fully describe individual dust amplitudes (e.g., the estimate is off by about 20% for on LR33). The scaling can help to asses the order of magnitude of the dust polarization level on a specific region, but is not a substitute for actually characterizing the polarized angular power spectra. Fig. 5 Ratio of the amplitudes of the and dust power spectra at 353 GHz for the different LR regions defined in Sect. 3.3.1 , distinguished here with . The mean value is plotted as a dashed line. 4.4. Amplitude of <inline-formula><alternatives><tex-math id="tex_eq104"></tex-math><inline-graphic mime-subtype="png" mimetype="image" xlink:href="aa25034-14-eq211.png"></inline-graphic><mml:math id="mml_eq104" type="inline"> <mml:msubsup> <mml:mi mathvariant="normal">D</mml:mi> <mml:mi mathvariant="italic">l</mml:mi> <mml:mi mathvariant="slanted">BB</mml:mi> </mml:msubsup> </mml:math></alternatives></inline-formula> relative to <inline-formula><alternatives><tex-math id="tex_eq105"></tex-math><inline-graphic mime-subtype="png" mimetype="image" xlink:href="aa25034-14-eq212.png"></inline-graphic><mml:math id="mml_eq105" type="inline"> <mml:msubsup> <mml:mi mathvariant="normal">D</mml:mi> <mml:mi mathvariant="italic">l</mml:mi> <mml:mi mathvariant="slanted">EE</mml:mi> </mml:msubsup> </mml:math></alternatives></inline-formula> We examine the ratio of the amplitudes of the fitted power laws found in Sect. 4.3 . The ratios
Paper: 2018A&A...615A.139A (acknowledgments)
Acknowledgements We thank the anonymous referee for their careful reading of and helpful suggestions on the manuscript, and Andreas Korn for providing the FOCES spectra used in this work. AMA and KL acknowledge funds from the Alexander von Humboldt Foundation in the framework of the Sofja Kovalevskaja Award endowed by the Federal Ministry of Education and Research, and KL also acknowledges funds from the Swedish Research Council (grant 2015-004153) and Marie Skłodowska Curie Actions (cofund project INCA 600398). TN and MA acknowledge funding from the Australian Research Council (grant DP150100250), and MA also acknowledges funding through ARC Laureate Fellowship (FL110100012). Parts of this research were conducted by the Australian Research Council Centre of Excellence for All Sky Astrophysics in 3 Dimensions (ASTRO 3D), through project number CE170100013. PSB acknowledges financial support from the Swedish Research Council and the project grant “The New Milky Way” from the Knut and Alice Wallenberg Foundation. Funding for the Stellar Astrophysics Centre is provided by The Danish National Research Foundation (grant DNRF106). This work made use of data from the UVES Paranal Observatory Project (ESO DDT Program ID 266.D-5655). This work was supported by computational resources provided by the Australian Governmentthrough the National Computational Infrastructure (NCI) under the National Computational Merit Allocation Scheme.
Paper: 2017A&A...606A..91S (fulltext)
XMM-Newton spectra of HD 54879 (high-resolution RGS spectra have an insufficient signal and are not useful). The X-ray spectra of magnetic hot stars are usually well described by multi-temperature thermal plasma models (e.g., Oskinova et al. 2011 ; Nazé et al. 2014 ). This is also the case for HD 54879. As a first step, we fitted the observed spectra in the 0.2–10.0 keV band with a thermal two-temperature spectral model that assumes optically thin plasma in collisional equilibrium. At this step, the absorption was modeled as originating in the cold interstellar medium ( Wilms et al. 2000 ). The resulting spectral fit is statistically significant, with a reduced . The best fit parameters are shown in Table 3 . There is no physical motivation for restricting the plasma temperature distribution to two temperatures. In fact, spectroscopic analyses show a continuous temperature distribution in hot stars ( Wojdowski Schulz 2005 ). Therefore, as a next step, we fitted a three-temperature spectral model to the observed EPIC spectra of HD 54879. Also in this case, the fit is statistically significant, with . The three-temperature model fit is shown in Fig. 10 , and the associated parameters are shown in Table 3 . The three-temperature fit reveals the hottest plasma component with a temperature of about 2 keV (20 MK). Better S/N is required to confirm the presence of this high-temperature plasma. We also tried to fit a four-temperature plasma model. However, the quality of the data is not sufficient to further constrain the temperature distribution. Fig. 10 XMM-Newton PN, MOS1, and MOS2 spectra of HD 54879 (black, red, and green curves, respectively) with error bars corresponding to 3 with the best fit three-temperature thermal model (solid lines). The model parameters are shown in Table 1. Table 3 X-ray spectral parameters derived from the XMM-Newton observations of HD 54879 assuming a three-temperature plasma model. Notes. (a) Corrected for the interstellar absorption in the 0.2–12 keV band. In the framework of the magnetically confined wind shocks (MCWS) model ( Babel Montmerle 1997 ; ud-Doula Owocki 2002 ), the wind plasma streams collide at the magnetic equator and give rise to a shock that heats the plasma. Hence, the maximum temperature follows from a Rankine-Hugoniot condition and cannot exceed a value determined by the maximum streaming velocity. Motivated by the recent discovery of a non-thermal component in the X-ray emission of the magnetic star HR 7355 ( Leto et al. 2017 ), we also fitted the observed spectra with a composite two-temperature thermal plasma and a non-thermal, power law component. The model provides a fit of similar quality to the observed spectra ( for 138 degrees of freedom). The best fit thermal plasma temperatures are [keV], [keV], and the power-law exponent is . Thus, the quality of the data is not sufficient to discriminate between three-temperature thermal and two-temperature thermal plus non-thermal component models.
Paper: 2019A&A...621A..40E (acknowledgments)
X-COP data products are available for download at https://www.astro.unige.ch/xcop . We thank Ricardo Herbonnet and Henk Hoekstra for providing us with their weak-lensing results in advance of publication, and the anonymous referee for helpful comments. The research leading to these results has received funding from the European Union’s Horizon 2020 Programme under the AHEAD project (grant agreement no. 654215). Based on observations obtained with XMM-Newton , an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA. S.E. acknowledges financial contribution from contracts NARO15 ASI-INAF I/037/12/0, ASI 2015-046-R.0 and ASI-INAF n.2017-14-H.0. F.V. acknowledges financial support from the ERC Starting Grant “MAGCOW”, no.714196. M.G. is supported by NASA through Einstein Postdoctoral Fellowship Award Number PF5-160137 issued by the Chandra X-ray Observatory Center, which is operated by the SAO for and on behalf of NASA under contract NAS8-03060. Support for this work was also provided by Chandra grant GO7-18121X. E.R acknowledges the ExaNeSt and Euro Exa projects, funded by the European Union’s Horizon 2020 research and innovation programme under grant agreements no. 671553 and no. 754337 and financial contribution from the agreement ASI-INAF n.2017-14-H.0. H.B. acknowledges financial support from ASI Grant 2016-24-H.0.
Paper: 2021MNRAS.500.1285B (fulltext)
negligibly, which suggests that for those cases, the lengthening of the spiral arms resulting from stripping can also give an appearance of unwinding without actually affecting the pitch angles. Schulz Struck ( 2001 ) noted possible winding of the inner spiral arms due to compression and ‘annealing’ of the disc in addition to unwinding of the outer regions of their face-on stripped galaxies. With the given sample, we do not have sufficient evidence to investigate this; however, with a larger sample of galaxies the effect of RPS on the more central spiral arms could be tested. The unwinding tails were found to be comprised of only the youngest stellar populations, in many cases only stars 20 Myr, when analysed with the sinopsis spectrophotometric fitting code. These have likely formed in situ , confirming that the unwinding effect only occurs in the gas component of the galaxy. Moreover, the age gradients observed in some unwinding arms possible capture the process in action, revealing a trail of progressively younger stars left by the opening arms of stripped material. Simulations were carried out of RPS interactions in different inclinations with respect to the ICM wind to investigate how the unwinding effect can arise and to compare the resulting morphology with observed features in the MUSE sample. The resulting snapshots were analysed throughout the stripping process, and used to produce ‘observations’ from different viewing angles to compare with the observed sample of galaxies. The two main effects that can, individually and in combination, give rise to an unwinding pattern are as follows: In edge-on wind interaction cases, the differential ram pressure in the leading and trailing edges of the disc causes a build-up of material in the leading edge. The tail becomes asymmetric and the slowed rotation of the gas causes the pitch angle to increase. In face-on wind interactions, the gas from the outer spiral arms is stripped and withdrawn from the potential well of the galaxy disc, causing it to move to higher orbits and increase the pitch angle in all directions from the disc. In most cases, a combination of these two effects will occur as a galaxy is stripped at angles between face-on and edge-on with respect to the ICM wind. Cases in which the galaxy is moving mostly edge-on to the wind appear to be more asymmetric and have tails that are further collimated to one side of the disc, while cases closer to face-on such as JO201 and JO85 show gas extending in a much wider range of directions from the disc. Furthermore, the simulations, combined with the phase-space positions of our galaxies and the sinopsis stellar population analysis, as well as similar effects found in the literature (Steinhauser et al. 2016 ; Roediger et al. 2014 ), all point towards the idea that the unwinding effect occurs only for galaxies undergoing first infall on radial plunging orbits, and appears to be a short-lived phase during stripping, containing stars
Paper: 2016A&A...591A..70K (acknowledgments)
This work was supported by the Swedish Research Council. M.M. has received funding from the People Programme (Marie Curie Actions) of the EUs FP7 (FP7/20072013) under REA grant agreement No. 623898.11. W.V. acknowledges support from ERC consolidator grant 614264.
Paper: 2018MNRAS.479.3498D (acknowledgments)
ACKNOWLEDGEMENTS The Torino polarimeter was built at the INAF (Instituto Nazionale Di Astrofisica) – Torino Astrophysical Observatory using dedicated funds provided in the framework of the INAF PRIN 2009 program. This work is based on data obtained at the C2PU (Centre Pédagique Planète et Univers) facility (Calern observing station, Observatoire de la Côte d’Azur, Nice, France) and at the Rozhen National Astronomical Observatory (Bulgaria). GB, ZD, TB, IB, AC, SB, and YK gratefully acknowledge observing grant support from the Institute of Astronomy and Rozhen National Astronomical Observatory, Bulgarian Academy of Sciences.
Paper: 2016MNRAS.456.4111B (fulltext)
115 ks with 25 ks on-source exposure. 4.1 <italic>INTEGRAL</italic> context Previous studies of this source with INTEGRAL have revealed a long (30 d), eccentric (e ∼ 0.3–0.4) orbit and shown detections in more than 60 per cent of neutron star periastron passages. SAX J1818.6−1703 was not detected by INTEGRAL in an observation carried out between 2013 March 12 and 2013 March 13 during Revolution 1271, at a time that corresponds to an orbital phase of 0.804–0.862. During this region of a long eccentric orbit, the stellar wind, through which the neutron star passes, is not expected to be sufficiently dense or structured so as to produce detectable emission from accretion. A bright flare from SAX J1818.6−1703 detected by Swift /BAT on 2013 March 17 (corresponding to phase, φ = 0.970) was reported by Romano et al. ( 2014 ). During this flare the source reached a peak flux of 146 mCrab (15–50 keV), clearly showing that the source exhibited enhanced X-ray emission a few days prior to the XMM observation reported here. The source was also detected by INTEGRAL during Revolution 1274 at a significance of ∼13σ. A more detailed analysis of this revolution shows that the source was first observed between UTC 11:10:03 and UTC 20:44:48 on 2013 March 20, corresponding to orbital phases φ = 0.086–0.099, with a total source exposure of 7.5 ks. During this observation, the source was detected with an average count rate of 2.17 ± 0.22 counts s −1 (18–60 keV) or ∼10 mCrab. The source was subsequently observed between UTC 09:00:11 and UTC 20:13:27 on 2013 March 22, corresponding to orbital phases φ = 0.150 – 0.165, again with a total source exposure of 7.5 ks. The average count rate in this observation was 1.27 ± 0.22 counts s −1 (18–60 keV). In order to compare these fluxes to previous periastron passages, we performed the recurrence analysis of Bird et al. ( 2008 ) using an INTEGRAL /IBIS light curve of SAX J1818.6−1703 from the first 1000 revolutions of INTEGRAL . We calculated the mean flux detected by INTEGRAL during a 4-day period centred on periastron and from this, we find that during the first observation of Revolution 1274, the source is brighter than 80 per cent of periastron passages observed by INTEGRAL . The flux detected between UTC 09:00:11 and UTC 20:13:27 on 2013 March 22 is consistent with fluxes observed by INTEGRAL during the majority of periastron passages. However, it is different from the distribution of apastron flux measurements and still represents a significant detection of the source. When considering both the Swift /BAT bright flare on 2013 March 17 and the INTEGRAL observations during Revolution 1274, it is clear that SAX J1818.6−1703 was in an atypically active state during the periastron passage covered by the INTEGRAl and XMM observations presented in this work. 4.2 Source distance Knowledge of the source distance is critical in order to calculate associated luminosities and hence disentangle the accretion processes
Paper: 2016A&A...589A..76G (acknowledgments)
This work has been supported by the Israeli-French ASTROPHYSICS program and by the French GdR PCHE/PNHE. Ehud Behar received funding from the European Unions Horizon 2020 research and innovation programme under the Marie Sklodowska-Curie grant agreement No. 655324, and from the I-CORE program of the Planning and Budgeting Committee (grant number 1937/12). The authors are gratetful to Suzy Collin, Jonathan Stern and Agata Różanská for their helpful comments on this paper. Finally, we thank the anonymous referee for constructive remarks.
Paper: 2021MNRAS.500.2461S (fulltext)
to orbital phase 0.9364, close to the primary eclipse, for which the contribution of the secondary is minimum. Figure 5. Same as Fig. 4 , but for one of the MagE spectra corresponding to orbital phase 0.9352. The results of the quantitative spectral analysis of all ALFOSC and MagE spectra are summarized in Fig. 6 . In the left-hand panels (ALFOSC), the apparent variations of the effective temperature with phase and an amplitude of ∼2500 K are obvious. The lowest temperatures (∼40 kK) occur near primary eclipse and in the secondary eclipse, where the contribution by extra light should be minimal. Hence, the increase during other orbital phases is caused by reflected light and, therefore, not real. Similarly, variations of the helium abundance are observed. The apparent variations of the surface gravity, however, are small. The mean values adopted for T eff , log g , and log y , summarized in Table 2 , are obtained by selecting six spectra close to the primary eclipse with phase between −0.1 and +0.1 and adding one spectrum at phase 0.52. The analysis of the MagE spectra also results in effective temperatures and helium abundances that seem to vary with orbital phase (cf. Fig. 6 , right-hand panels), but with amplitudes less pronounced than those from the ALFOSC spectra. For this reason, we use all the MagE spectra to compute mean values and standard deviations of T eff and log y (cf. Table 2 ). Figure 6. T eff , log g , log y , and v rot sin i variations as a function of the orbital phase. Left-hand panels: ALFOSC. The green dotted–dashed horizontal lines mark the adopted T eff , log g , and log y and the associated errors (cf. Table 2 and text). Note the two points with low T eff near phase 0.5 (secondary eclipse), when the contribution of the secondary is strongly reduced. Right-hand panels: MagE. The orange dotted–dashed horizontal lines mark the average values and associated errors of T eff , v rot sin i , and He abundance. In the central panel, note the outlier near phase zero (primary eclipse). Table 2. SdOB atmospheric parameters and rotational velocity. The effective temperature and surface gravity of EPIC 216747137 ( T eff = 40400 K and log g = 5.56) are very similar to the hot HW Vir systems AA Dor (Klepp Rauch 2011 ) and V1828 Aql (= NSVS 14256825, Almeida et al. 2012 ). EPIC 216747137 also shares an underabundance of helium (log y = −2.59) with the two others. The projected rotational velocity, as derived from the individual MagE spectra (central right panel of Fig. 6 ), results in a mean v sin i = 51 ± 10 km s −1 , significantly less than ∼70 km s −1 expected for tidally locked rotation. Mean rotational velocity of 51 km s −1 and standard deviation of 10 km s −1 are obtained excluding only a single outlier close to phase 0 (see central right panel of Fig. 6 ). 5 STELLAR PARAMETERS: RADIUS, MASS, AND LUMINOSITY The second data release of Gaia provided a precise (5 per cent) parallax measurement which allows the stellar
Paper: 2016ApJ...832...11B (acknowledgments)
Dedicated to the memory of Janet C. Bradford . We thank the anonymous referee and Stacy McGaugh for helpful comments that have improved the quality of this manuscript. We would like to thank Stephane Courteau, Susan Kassin, Allison Merritt, and Erik Tollerud for helpful feedback on this manuscript, as well as Michael Blanton and the BS group for helpful discussions. J.D.B. acknowledges support from the Gruber Foundation and the National Science Foundation Graduate Research Fellowship Program. M.G. acknowledges a fellowship from the John S. Guggenheim Memorial Foundation. The Arecibo Observatory is operated by SRI International under a cooperative agreement with the National Science Foundation (AST-1100968), and in alliance with Ana G. Méndez-Universidad Metropolitana, and the Universities Space Research Association. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. Funding for SDSS-III has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, and the U.S. Department of Energy Office of Science. The SDSS-III web site is http://www.sdss3.org/ . This material is based upon work supported by the National Science Foundation Graduate Research Fellowship Program under Grant No. DGE-1122492. Any opinions, findings, and conclusions or recommendations expressed in this material are those of the author(s) and do not necessarily reflect the views of the National Science Foundation.
Paper: 2019A&A...629A..93Y (fulltext)
(2009) applied an early version of the IGIMF theory where the IMF variation of individual star clusters was not applied. The modern IGIMF prescription used here with a possible top-heavy gwIMF further increases the resulting [Mg/Fe] value. 7. Conclusions This contribution publishes an open source galaxy evolution code that is able to vary the gwIMF according to the IGIMF theory (or any other variable IMF theory) at each time step, yielding the mass composition, supernova rate, and chemical composition evolution of the galaxy. Following Yan et al. (2017) and Jeřábková et al. (2018) , where the gwIMF for a single star formation epoch is derived according to the IGIMF theory, the newly developed code published here enables us, for the first time, to perform a comprehensive and detailed examination of the effect of applying a different IMF theory in the context of galaxy chemical evolution assuming an identical SFH. Our preliminary result of a high-mass starburst galaxy assuming different IMFs and the same SFH is shown in Sect. 6 ; the result assuming the same gwIMF (the IGIMF theory) but different SFHs is shown in Appendix A . A systematically varying gwIMF leads to a strong modification of the element composition of a galaxy (naturally reproducing the observed higher metallicity and α -enhancement of the massive elliptical galaxies), while the difference in SFH and recent star formation activity influences the interpreted property of a galaxy due to a difference in the luminosity-weighted observation. The star formation environment of massive ellipticals in the early Universe is very different from the local Universe where we calibrate the IGIMF theory. There is no guarantee that such an extrapolation of the theory would work, but it is certainly an inspiring step. We will use this open source code for more comprehensive and detailed studies for galaxies with different masses and SFHs and also develop a spatially resolved IGIMF formulation. Appendix A: Results for a log-normal SFH Fig. A.1. Same as Fig. 5 , but for the log-normal SFH shown in Fig. 2 that extends for 13 Gyr. Fig. A.2. Same as Fig. 6 , but comparing different SFHs instead of different IMF assumptions. The orange lines apply the boxy SFH and the red lines apply the log-normal SFH, both shown in Fig. 2 . The same colour-coding is applied from Figs. A.2 – A.7 . The serrated shape is caused by the dying stars with a lifetime shorter than 10 Myr (see Sect. 6.2 ). We note that the orange lines are the same as the orange lines in Fig. 6 . This section compares the simulation results applying the IGIMF theory, but with the two different SFHs shown in Fig. 2 . The log-normal SFH forms the same stellar mass, but on a more extended timescale. The tail of the SFH with a low SFR leads to a stellar population that receives a contribution every δ t = 10 Myr from a top-light gwIMF, which does not exist in the box-shaped SFH model described in the main text. Thus, a more realistic and extended