adsabs/WIESP2022-NER huggingface dataset:
Paper: 2020A&A...638L...4G (fulltext)
Megeath et al. 2012 ) and included in the Herschel Orion protostar survey (HOPS; Furlan et al. 2016 ) as source 383. The submillimeter to bolometric luminosity ratio of HOPS 383 , 1.4% ( Safron et al. 2015 ) well surpasses the (0.5%) threshold for bona fide Class 0 designation ( André et al. 1993 , 2000 ). Notably, HOPS 383 is the first Class 0 protostar known to have undergone a mass-accretion-driven eruption ( Safron et al. 2015 ), which peaked by 2008 and ended by 2017 September (Appendix B ). 2. Observations We observed HOPS 383 three times with the Chandra X-ray Observatory ( Weisskopf et al. 2002 ) from 2017 December 13 to 14 with simultaneous near-infrared imaging on 2017 December 14 using the 4.1 m Southern Astrophysical Research (SOAR) telescope ( Krabbendam et al. 2004 ) in Chile. We used Chandra ’s Advanced CCD Imaging Spectrometer (ACIS-I) in the very faint mode with a frame time of 3.141 s (Appendix C.1 ). Chandra ACIS-I has on-axis a full-width half-maximum (FWHM) angular resolution of and a FWHM spectral resolution of 261 eV at 6.4 keV. Data reduction is described in Appendix C.2 . The SOAR Spartan infrared camera ( Loh et al. 2012 ) is composed of 4 CCDs of 2048 × 2048 pixels. We selected the wide-field mode with a single-detector field-of-view of (i.e., edge-to-edge) and a pixel scale of with the K -band and H 2 narrow-band filters. We used a five-point dithering pattern where, in the first exposure, HOPS 383 was put near the center of the northeastern detector (det3), which was then moved by 30′′ from its initial position sequentially toward the south, north, west, and east in the four other exposures. This sequence was repeated until the requested exposure was achieved (Appendix D.1 ). Data reduction is described in Appendix D.2 . 3. Near-infrared results The near-infrared nebulosity that was prominent during the outburst ( Safron et al. 2015 ) is not detected in our SOAR K -band image (Fig. 1 a), indeed, it vanished by 2015 December 30 ( Fischer Hillenbrand 2017 ). The H 2 -narrow filter image shows an emission knot of shocked molecular hydrogen at ∼15′′ southeast from the mid-infrared location of HOPS 383. We identify this H 2 source with the H 2 emission knot SMZ 1-2B ( Stanke et al. 2002 ) that was observed on 1997 September 13, yielding a proper motion of southeast in 20.25 yr (Fig. 1 b). The corresponding velocity is ∼(95 ± 53)×(sin i /sin69.5°) −1 × ( d /420 pc) km s −1 , where d is the distance to HOPS 383 and i is the inclination angle (90° corresponds to an edge-on view; Table 1 of Furlan et al. 2016 ). This is typical of the velocity values observed in the outflows from Class 0 protostars ( Bontemps et al. 1996 ; Reipurth Bally 2001 ). Assuming a constant velocity, the angular distance of ∼16′′ between HOPS 383 and this emission knot corresponds to a kinematical timescale of ∼180 ± 100 yr. The proper motion direction is consistent with the orientation of the outburst nebula, the 12 CO blueshifted outflow (
Paper: 2021MNRAS.500..772D (fulltext)
by the electrons in the converging flow on to a compact star and they called such Comptonization BMC, i.e. Bulk-Motion Comptonization. The BMC model has been adopted to interpret the PL component in the broad-band X-ray spectrum of BHXBs well (Shrader Titarchuk 1998 ; Titarchuk Seifina 2009 ; Titarchuk Shaposhnikov 2010 ). The BMC is one of the models of xspec and it is formulated as This expression consists of two subcomponents: one is the injected blackbody-like component ( BB ) to provide the seed photons for Comptonization, and the other is BB * G to describe the Comptonization emission, where G is the response of the Compton medium. In xspec , this model includes four parameters, i.e. kT bb , α, log ( A ), and C N , which are the seed photon temperature, spectral index (photon index Γ = α + 1), logarithm of weighting factor, and normalization, respectively. The observed fluxes of BMC are comprised of a direct emission flux and a Comptonized emission flux. The fraction of the direct emission is calculated with the factor 1/(1 + A ), whereas the fraction of Comptonized emission is inferred by a Comptonization factor f ( f = A /(1 + A )). In two cases, the BMC model is reduced to BB: one is with A ≪ 1, indicating that the value of f is very small, so the Comptonized emission is negligible; the other is with α ≫ 1, meaning that the energy for the seed photons to get from the high-energy electrons is very small, so the Comptonization is inefficient. The theoretical study of Titarchuk et al. ( 1997 ) shows that the BMC process taking place around the compact star can only result in the PL component with Γ > 1. Therefore, theoretically, the BMC model can only account for the relatively soft hard X-ray PL tail in the X-ray spectrum of XRBs. In xspec , in order to ensure Γ > 1, the spectral index α is set to be a positive parameter, i.e. α > 0. In this work, the obtained photon index of the PL hard tail spans a range of ∼(1.1–2.3) and ∼(0.7–2.1) on the HB and on the NB, respectively. Considering the error of the lower limit of NB photon index, the photon index of HB or NB satisfies the condition of Γ > 1. Assuming the hard X-ray tail in Sco X-1 to be originated from the BMC process around the NS, we form model to fit the broad-band spectra of this source. Assuming the NS to provide the seed photons for the dynamic Comptonization in the converging accretion flow on to the compact star, the BB emission of the NS as well as the outgoing PL component is included in the BMC model. Adopting that there exists an ADC in Sco X-1, the disc emission is described by a CPL (Church et al. 2012 ). Taking into consideration the iron line and the additive BB in the PCA spectrum, finally, we form a multicomponent model of BMC+LINE+CPL+BB to fit the PCA+HEXTE spectra of HB and NB in which hard X-ray PL tails are detected. The fitting results are listed in Table 4 and four unfolded spectra are shown in Fig. 9 . As listed in Table 4 , the values of the reduced
Paper: 2018A&A...610A..81L (acknowledgments)
Acknowledgments The author is grateful to an anonymous referee for a careful reading of the manuscript and several suggestions that helped him to substantially improve the present work. AFL wishes to thank Drs. A. Maggio, S. Sciortino, I. Pillitteri, I. Pagano, and S. Desidera for interesting discussions on magnetic star–planet interactions. This work was partially supported by the project WOW, one of the Progetti Premiali of the Italian Ministry of Education, University, and Research funded to the Italian National Institute for Astrophysics (INAF).
Paper: 2015A&A...578A.135V (acknowledgments)
We thank the staffs of the ESA Herschel Science Centre and NASA Herschel Science Center, and the HIFI, Instrument Control Centre (ICC) for their help with the data reduction routines. In addition, we owe special thanks to David Teyssier for clarifications regarding the hebCorrection tool. This work was performed at the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.
Paper: 2020MNRAS.493...78A (acknowledgments)
ACKNOWLEDGEMENTS This work is partly supported by Ministry of Economy and Competitiveness (MINECO) (Spain), FIS2016-76363-P, by COST Action PHAROS (CA16214), by project 2017 SGR247 Agency for Management of University and Research Grants (AGAUR, Catalonia) (SDO). AVA thanks the Program 5–100 Immanuel Kant Baltic Federal University (IKBFU, Russia).
Paper: 2017MNRAS.471.4345Q (fulltext)
= 5 in the top panel of Fig. 5 . The result using the Tiamat halo merger trees is shown for comparison and to indicate the convergence of the quasar luminosity function at z ∼ 5. The model is in agreement with observations at the bright end (Shen Kelly 2012 ; McGreer et al. 2013 ). We indicate the quasar luminosities of the three analogues in the top panel of Fig. 5 and find that the progenitors of ZF-1, ZF-2 and ZF-3 are bright quasars in the range that optical surveys such as SDSS are able to detect. This suggests that ZF-COSMOS-20115 is likely the descendent of a high-redshift quasar and illustrates the possible connection between MQGs and high-redshift quasars (Maiolino et al. 2012 ). Figure 5. Top panel: quasar UV 1450 Å luminosity function at z ∼ 5. The results using the Tiamat-125-HR and Tiamat halo merger trees are shown with solid and dashed lines, respectively. The shaded regions represent the 95 per cent confidence intervals around the mean using 100 000 bootstrap re-samples for the Tiamat-125-HR result. The observational data is shown with different symbols: Shen Kelly ( 2012 ) using SDSS DR7 data at z ∼ 4.75 () and McGreer et al. ( 2013 ) using SDSS, UKIDSS and MMT at z ∼ 4.7–5.1 (▵). From left to right, the vertical lines indicate the quasar UV magnitudes of ZF-1, ZF-2 and ZF-3, respectively. Bottom panel: correlation between high-redshift black hole mass ( z = 5) and low-redshift SFR ( z = 3.7). The size of the empty circle represents different stellar masses and ZF-3 is indicated with the filled circle. In order to better demonstrate this, for each massive galaxy selected at z = 5.3, we present the correlation between black hole mass and subsequent SFR at z = 3.7 in the bottom panel of Fig. 5 . ZF-3 is indicated with a filled circle. We see that galaxies with less massive black holes ( M BH 10 7.5 M ⊙ ) can have a range of subsequent star formation levels. This is due to self-regulated stellar feedback being the dominant mechanism in these galaxies. However, due to persistent AGN feedback, galaxies with more massive black holes do not have intense subsequent star formation. This indicates that a high-redshift bright quasar will likely become a MQG at later times, when its accretion disc has been completely consumed. 5.3 The reionization history The three ZF-COSMOS-20115 analogues are hosted by massive haloes, which provide deep gravitational potentials that efficiently accrete baryons. The high baryon accretion efficiency induces intense star formation at high redshift, providing a significant number of ionizing photons that reionize surrounding hydrogen at a very early time. In this section, we discuss the reionization history of ZF-COSMOS-20115 . In Fig. 6 , we show the evolution of the average neutral hydrogen fraction within boxes with different sizes around the three analogues. The neutral hydrogen fraction is calculated using 21 cmFAST (Mesinger, Furlanetto Cen 2011 ) with the galaxy catalogue provided by Meraxes (see
Paper: 2020MNRAS.498.4316B (fulltext)
from (numerical) frequency runs. In that case, reorganization energies (e.g. ) computed from single-point calculations were utilized to get the adiabatic values. APPENDIX A2: ADDITIONAL DETAILS ON THE CALCULATION OF ROVIBRATIONAL PROPERTIES Deceptive as it may sound for a theoretical chemist, like in the case of other carbon chains (e.g. Bâldea 2020b ), the DFT approach appears to perform better in quantifying rovibrational spectra for MgC 4 H species than much more computationally demanding ab initio methods: For MgC 4 H 0 neutral, our very tight optimization at B3LYP/aug-cc-pVTZ level yielded MHz. The relative deviation from the experimental value MHz (Cernicharo et al. 2019 ) is per cent, which is significantly better than both 0.3 per cent at CCSD(T)-F12/cc-pCVTZ-F12 (Cernicharo et al. 2019 ) and 1.7 per cent at CCSD(T)/aug-cc-pVTZ (Cernicharo et al. 2019 ). Likewise, our value D DFT = 0.0641 kHz obtained at the B3LYP/aug-cc-pVTZ level is only 13 per cent smaller than the experimental value D exp = 0.074 ± 0.002 kHz (Cernicharo et al. 2019 ). This is better than 17 per cent, the deviation of CCSD(T)-value (Cernicharo et al. 2019 , and Table A2 ), which, in turn, is somewhat better than 18 per cent, which is the deviation of the CCSD(T)-F12-based value (Cernicharo et al. 2019 ). Noteworthily, the best estimate D Kratzer = 0.0741 kHz (cf. Table 4 and A3 ) is that based on the Kratzer formula (equation 1 ) used in conjunction with the B3LYP/aug-cc-pVTZ B and ω MgC . This Kratzer-based value cannot be distinguished from the experimental value D exp = 0.074 ± 0.002 kHz (Cernicharo et al. 2019 ). Table A3. Calculated centrifugal distortion constants D (kHz) presented along with estimates deduced via Kratzer approximation (equation 1 ) using harmonic frequencies ω MgC of the MgC stretching mode. a CCSD(T)/cc-pVTZ. b CCSD(T)/aug-cc-pVTZ. c B3LYP/6-311++G(3DF, 3PD). d B3LYP/aug-cc-pVTZ. e CCSD(T)/cc-pCVTZ. Table A4. Adiabatic and vertical values of electron attachment (EA) and ionization (IP) energies computed for MgC 4 H linear chains at several levels of theory. Their most reliable estimates are depicted in italics and were used to deduce the chemical hardness η, chemical softness S , chemical potential μ, Mulliken electronegativity χ, and electrophilicity index ω. Energy unit is eV. a Δ-CCSD(T)/cc-pVTZ. b EOM-CCSD/cc-pVTZ. c Δ-CCSD(T)/cc-pCVTZ. d EOM-CCSD/cc-pCVTZ. e Δ-CCSD(T)/aug-cc-pVTZ. f .EOM-CCSD/aug-cc-pVTZ. g Δ-B3LYP/6-311++G(3DF, 3PD). h CBS-QB3. For sure, neither a DFT/B3LYP/aug-cc-pVTZ calculation like that found performing the best above nor whatever other much more sophisticated theoretical approach currently available is precise enough to enable detection in space. In view of this state-of-the-art of theory, aiming at guiding subsequent laboratory measurements and interpretation of not yet available rovibrational spectra (as the case of ), it is current practice in astrophysics (e.g. Pauzat et al. 1991 ) to estimate rotational
Paper: 2018ApJ...867..156V (acknowledgments)
This work was supported by NSF SHINE grant AGS1622413 to the University of New Hampshire (UNH). C.W.S. was also supported by Caltech subcontract 44A-1062037 to UNH in support of the ACE/MAG instrument. J.E.S. is supported by STFC(UK) grant ST/N000692/1.
Paper: 2018ApJ...858....8N (fulltext)
origin ( t = 61 minutes) up to its decay phase ( t = 70.7 minutes) in the vertical experiment for the three LOSs. (An animation of this figure is available.) To extend the analysis, it is also of interest to consider the total emitted intensity for each vertical column, i.e.,which, following Equation ( 1 ), is equal to the column density of the emitting species along the LOS (except for a constant factor). Equation ( 5 ) has been calculated separately for Si iv and O iv and with the emissivities obtained assuming either NEQ or SE ionization, to better gauge the importance of disregarding the NEQ effects. The results are shown in the middle and bottom panels of Figure 10 for Si iv and O iv , respectively. The top panel of the image contains the 2D map of the emissivity for Si iv for context identification. Combining Figures 9 and 10 , we are able to discern and describe characteristic features of the spectral profiles: 1. The most prominent feature in the synthetic profiles of Figure 9 is the brightening associated with the location of the internal footpoint of the surge. In the corresponding movie, we can see how that footpoint is formed at around t = 64 minutes, as the surge detaches from the emerged dome. During those instants, the associated synthetic profiles are characterized by large intermittent intensities and bidirectional behavior with velocities of tens of kilometers per second, as apparent, e.g., in the (b0) and (c0) panels at x ∈ [15, 16] Mm. In Si iv and O iv (panels (b) and (c)), we find that the internal footpoint is usually the brightest region, although there are some instants in which the brightest points can be located in the crests or the external footpoint. This is a potentially important result from the observational point of view because it can help us to unravel the spatial geometry of the surge in future observational studies: if strong brightenings are detected in Si iv and also in O iv within the surge, it would be reasonable to think that they correspond to the internal footpoint of the surge. In this region, the intensity ratio between Si iv and O iv ranges between 2 and 7, approximately. Note that, in general, the intensity ratio values vary depending on the observed region and features (Martínez-Sykora et al. 2016 ). In Section 5.2 we will see that LOS effects play a major role in causing the large brightness of the internal footpoint (and other bright features) compared to the rest of the surge. Here we consider the parallel question of the role of NEQ: what would be obtained for the intensity of the internal footpoint if SE were assumed? Comparing the values for Si iv in Figure 10 (middle panel, 13.8 ≤ x ≤ 16.1), there is roughly a factor of two, on average, between the NEQ and SE intensities; for O iv (bottom panel), there is no major difference in the intensity between both calculations. One could conclude that while NEQ is important for the Si iv diagnostics, SE could be applied for the O iv case;
Paper: 2015ApJ...810...25G (acknowledgments)
The Pan-STARRS1 Surveys (PS1) have been made possible through contributions of the Institute for Astronomy, the University of Hawaii, the Pan-STARRS Project Office, the Max-Planck Society and its participating institutes, the Max Planck Institute for Astronomy, Heidelberg, and the Max Planck Institute for Extraterrestrial Physics, Garching, The Johns Hopkins University, Durham University, the University of Edinburgh, Queen's University Belfast, the Harvard-Smithsonian Center for Astrophysics, the Las Cumbres Observatory Global Telescope Network Incorporated, the National Central University of Taiwan, the Space Telescope Science Institute, the National Aeronautics and Space Administration under Grant No. NNX08AR22G issued through the Planetary Science Division of the NASA Science Mission Directorate, the National Science Foundation under grant No. AST-1238877, the University of Maryland, and Eotvos Lorand University (ELTE). This publication makes use of data products from the 2MASS, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation. The computations in this paper were run on the Odyssey cluster supported by the FAS Science Division Research Computing Group at Harvard University. G.G. and D.F. are supported by NSF grant AST-1312891. E.S. acknowledges funding by Sonderforschungsbereich SFB 881 “The Milky Way System” (subproject A3) of the German Research Foundation (DFG). N.M. gratefully acknowledges the CNRS for support through PICS project PICS06183.
Paper: 2018MNRAS.476.4272E (acknowledgments)
ACKNOWLEDGEMENTS We thank the anonymous referee for helpful comments that improved this paper. We gratefully acknowledge support from the Science and Technology Facilities Council (HPE through grant ST/K501979/1, SR through ST/N50404X/1, TPR through ST/P000541/1). HPE acknowledges support under NASA contract NNG08FD60C. This research made use of data obtained from the 3XMM XMM–Newton Serendipitous Source Catalogue, based on archival observations with XMM–Newton , an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA. We also acknowledge the use of public data obtained from the Chandra Data Archive and the Swift Data Archive.
Paper: 2019A&A...621A.116H (fulltext)
which may not be particularly informative, as Fig. 9 demonstrates. To this end, we do not interpret the total flux from the nebula in terms of any particular spectral model; rather, we present a total integrated flux of the nebula emission to interpret the overall source strength. The resulting spectral flux obtained for the nebula from the summation of the 38 independent 0.26°× 0.26° spectral boxes is shown in Fig. 10 . For illustration, the HESS J1825–137 spectrum obtained from Fermi -LAT ( Grondin et al. 2011 ) is also shown, although we recall that the region from which this spectrum is obtained, based on a morphological symmetric Gaussian fit template to the Fermi data, does not fully correspond to the same spatial region. The 2011 analysis and the 3FGL (the third LAT Catalogue, 20 MeV to 300 GeV, Acero et al. 2015 ) both use a Gaussian with σ = 0.56° as the spatial template, whereas the 3FHL (the third LAT High-Energy Catalogue, 10 GeV to 2 TeV, Ajello et al. 2017 ) uses a Gaussian with σ = 0.75°, which is consistent with the Fermi Galactic Extended Source catalogue best-fit Gaussian with σ = 0.79 ± 0.04 ( Ackermann et al. 2017 ). Whilst these templates are only approximations to the source morphology, in Fig. 10 the IC peak may be inferred to occur at around 200 GeV. Fig. 10 Spectral energy distribution of HESS J1825–137 showing the total combined H.E.S.S. energy flux with statistical error bars, as determined with the method outlined in Sect. A . The Fermi -LAT spectral fit from Grondin et al. (2011) is also shown, together with spectra from the 3FGL and 3FHL catalogues ( Acero et al. 2015 ; Ajello et al. 2017 ), both of which were best fit by a log-parabola spectral model. A plausible IC emission model with multiple contributing radiation fields is shown, using the parameterisation of the Popescu et al. (2017) radiation model and the Naima modelling package ( Zabalza 2015 ). The emission is mostly dominated by the FIR field; the CMB becomes dominant at a few tens of TeV and below ~ 10 GeV. The total flux above 300 GeV was found to be (6.07 ± 0.13 stat ± 0.07 sys ) × 10 −11 cm −2 s −1 , corresponding to ~54% of the flux of the Crab nebula. Above 1 TeV, a flux of (1.12 ± 0.03 stat ± 0.25 sys ) × 10 −11 cm −2 s −1 was obtained, such that the proportion relative to the Crab nebula increases slightly to ~ 64% of the Crab nebula flux. The source remains strong at ~61% of the Crab nebula flux above 10 TeV. It is impressive that the γ -ray flux is so strong given that the distance to HESS J1825–137 is approximately twice that of the Crab nebula (at 2 kpc), despite (or perhaps because of) the significantly older age of the system. Interestingly, as part of the second HAWC catalogue, it was found that the confused source associated with the combined emission from the region, 2HWC J1825–134, is brighter than the Crab nebula at 7 TeV ( Abeysekara et al. 2017 ). This point is included in Fig. 10 and is compatible within the HAWC systematic
Paper: 2016MNRAS.459.2525S (acknowledgments)
MJ-H is supported in this work by a Marsden Grant administered by the Royal Society of New Zealand. SS was supported by a Victoria Doctoral Scholarship. We would like to thank Dr Phil Edwards for allocating Director's time for the H214 ARDES observations. Additionally, SS deeply thanks Dr Robin Wark for her assistance with remote-observing, and Dr Jamie Stevens and Dr Mark Wieringa for their invaluable comments on CABB data reduction and imaging in MIRIAD . Finally, we thank Dr Siamak Dehghan for assistance with the generation of the multipanel figures. The ATCA is part of the Australia Telescope National Facility which is funded by the Australian Government for operation as a National Facility managed by CSIRO. The Digitized Sky Surveys were produced at the Space Telescope Science Institute under US Government grant NAG W-2166.