adsabs/WIESP2022-NER huggingface dataset:
Paper: 2016ApJ...829L..31D (fulltext)
2013 ) and radial velocities (Saar Donahue 1997 ; Korhonen et al. 2015 ). Stellar activity has even been the culprit behind false detections of exoplanets around M dwarfs in high-precision data (Robertson et al. 2014 ). The role of M dwarf magnetic activity on planetary habitability, particularly from X-ray and UV flux due to quiescent emission and flares, is an ongoing area of research (Scalo et al. 2007 ; Tarter et al. 2007 ; Seager Deming 2010 ). Planetary habitability is impacted by the properties and evolution of both the star and planet (e.g., Vidotto et al. 2013 ). While models from Segura et al. ( 2010 ) show that single large M dwarf flares do not pose a great risk to planetary habitability, young planets may be bombarded with much stronger and more frequent flares than previously predicted (Armstrong et al. 2016 ). M dwarf activity can strongly affect planetary atmospheres through photoevaporation on Gyr timescales (Cuntz Guinan 2016 ; Owen Mohanty 2016 ). Whether this has a net positive or negative impact on atmosphere loss and planetary habitability is debated (Luger et al. 2015 ; Owen Mohanty 2016 ). Accurate studies of M dwarf magnetic activity in X-rays, UV, and the blue optical are challenging due to the intrinsic faintness of these cool stars at short wavelengths. However, the nearest M dwarf to the Sun, Proxima Cen, has long been known as an active flare star (Thackeray 1950 ), with a moderately strong magnetic field for its spectral type of M5.5 (Reiners Basri 2008 ). At a distance of 1.3 pc, Proxima Cen enables unique characterization of activity in X-ray through optical (Güdel et al. 2004 ; Fuhrmeister et al. 2011 ), and even radio frequencies (Lim et al. 1996 ). High time resolution flare data from Proxima Cen has been used to better understand flare heating mechanisms (Reale et al. 2004 ; Kowalski et al. 2016 ). Decades-long spectroscopic monitoring has even revealed a possible activity cycle for Proxima Cen (Cincunegui et al. 2007 ). Proxima Cen has been the subject of multiple searches for exoplanets (e.g., Benedict et al. 1999 ; Endl Kürster 2008 ). Recently, Anglada-Escudé et al. ( 2016 ) discovered the existence of Proxima b, a 1.27 M ⊕ planet orbiting within the star’s habitable zone, using multiple years of high-precision radial velocity monitoring. Proxima Cen is therefore an important benchmark object for understanding planet formation around low-mass stars, the evolution of the magnetic dynamo, and the impact of stellar activity on planetary atmospheres. Using data from the Microvariability and Oscillations of STars microsatellite (hereafter MOST ; Walker et al. 2003 ), we have conducted a study of the white-light flares from Proxima Cen. Flares have previously been studied with MOST for the famous active M3.5 dwarf, AD Leo (Hunt-Walker et al. 2012 ). Despite its long history of study as an active M dwarf, the census of flares from Proxima has not been constrained by modern space-based photometric studies.
Paper: 2019A&A...630A..93M (fulltext)
at epoch t , the catalog positions at t 0 would be corrected by applying the Galactic aberration model correction. We have investigated what is the effect of the aberration on estimated source positions. Figures 3 – 6 show the Calc/Solve differences in source positions (RA, Dec) versus RA and Dec for an aberration constant A G of 5.8 μ as yr −1 , which is the value we recommended for the IAU ICR3 solution as discussed below in Sect. 6 . The variation (scatter) of the differences at RA or Dec in these plots is due to the fact that the difference in the mean epochs from the chosen reference epoch 2015.0 vary significantly over the set of sources. The source position differences range over ±30–40 μ as. Fig. 3. Change in right ascension due to aberration versus right ascension with an aberration constant of 5.8 μ as yr −1 and reference epoch of 2015.0. Fig. 4. Change in right ascension due to aberration versus declination with an aberration constant of 5.8 μ as yr −1 and reference epoch of 2015.0. Fig. 5. Change in declination due to aberration versus right ascension with an aberration constant of 5.8 μ as yr −1 and reference epoch of 2015.0. Fig. 6. Change in declination due to aberration versus declination with an aberration constant of 5.8 μ as yr −1 and a reference epoch of 2015.0. 4. Estimates derived from Galactic astronomy measurements Aberration can also be derived from recent (2009–2017) Galactic astronomy measurements (e.g., Reid et al. 2014 ; Rastorguev et al. 2017 ; Brunthaler et al. 2011 ). These measurements are trigonometric parallaxes and proper motions of masers in high-mass star-forming regions in the Milky Way galaxy. The measurements were made using the Very Long Baseline Array (VLBA), the European VLBI network (EVN), and the Japanese VLBI Exploration of Radio Astronomy Project (VERA). The most recent investigation noted here ( Rastorguev et al. 2017 ), used a maser sample of 136 sources. Using these parallax and proper motion measurements, different investigators have derived models of the galaxy. Among the parameters of these models are the radial distance R 0 (kpc) to the Galactic center and circular rotation speed Θ 0 (km s −1 ) of the solar system barycenter around the Galactic center, which is located at a position with right ascension and declination of α G = 266.4° and δ G = −29.0° ( Reid Brunthaler 2004 ). To determine an aberration constant, we just consider the effect of the Local Standard of Rest (LSR) circular revolution with speed Θ 0 around the Galactic center. We note that there could be a small pecular acceleration an order of magnitude less than galactocentric acceleration due to solar motion with respect to the LSR as discussed in the introduction. Based on the estimated parameters R 0 and Θ 0 and their uncertainties from each investigator, one can determine the acceleration vector amplitude of (see Kovalevsky 2003 ) and its uncertainty assuming a circular orbit. This can be converted to an aberration
Paper: 2016MNRAS.458.3927B (acknowledgments)
ZZ acknowledges support by NASA through Hubble Fellowship grant HST-HF-51333.01-A awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA, under contract NAS 5-26555. We thank Olivier Berné, Simon Casassus and Zsolt Regály for helpful discussions, and Eugene Chiang, the reviewer, for a prompt and constructive report.
Paper: 2015MNRAS.454.2423A (acknowledgments)
It is a pleasure to thank A. Grenzebach, Y. Mizuno, Z. Younsi and A. Zhidenko for useful discussions and comments. We are also grateful to the referee, D. Psaltis, for comments and suggestions that have improved the presentation. This research was partially supported by the Volkswagen Stiftung (Grant 86 866) and by the ERC Synergy Grant ‘BlackHoleCam – Imaging the Event Horizon of Black Holes’ (Grant 610058). AAA and BJA are also supported in part by the project F2-FA-F113 of the UzAS and by the ICTP through the projects OEA-NET-76, OEA-PRJ-29. AAA and BJA thank the Institut für Theoretische Physik Goethe Universität' for warm hospitality during their stay in Frankfurt.
Paper: 2020A&A...633A..83B (acknowledgments)
We would like to thank Tobias Wiengarten for assistance with the magnetic field data, as well as Malina Reytemeier, Paul Evenson and the entire IceCube collaboration, as well as members of the Ruhr Astroparticle Plasma Physics (RAPP) Center for helpful discussions on the subject. We are further grateful for the support from the BMBF (05A14PC1) and from MERCUR (St-2014-0040). This work utilizes data from the National Solar Observatory Integrated Synoptic Program, which is operated by the Association of Universities for Research in Astronomy, under a cooperative agreement with the National Science Foundation and with additional financial support from the National Oceanic and Atmospheric Administration, the National Aeronautics and Space Administration, and the United States Air Force. The GONG network of instruments is hosted by the Big Bear Solar Observatory, High Altitude Observatory, Learmonth Solar Observatory, Udaipur Solar Observatory, Instituto de Astrofísica de Canarias, and Cerro Tololo Interamerican Observatory.
Paper: 2019MNRAS.484.1506L (acknowledgments)
ACKNOWLEDGEMENTS We thank the referee for a thoughtful read of the paper and comments that helped strengthen our paper. We thank Alexander Stephan and Bao-Minh Hoang for useful discussion and suggestions. CL acknowledges the UCLA Van Tree Foundation, Undergraduate Research Scholars Program Scholarship as well as the Stone Foundation Summer Research Scholarship. CL thanks Izuki Matsuba for discussion and help. SN acknowledges the partial support of the Sloan fellowship and from the NSF through grant No. AST-1739160.
Paper: 2017ApJ...846L..26W (fulltext)
converged when the emcee walkers had reached a steady state with measured best-fit values changing minimally over a large number of steps. 3. Results Figure 1 shows the 1.3 mm dust continuum, the 12 CO (2–1) integrated intensity map (moment 0), and the intensity-weighted velocity map (moment 1) of the FW Tau C disk. Similar to Kraus et al. ( 2015 ) and Caceres et al. ( 2015 ), no signal is found from the close binary FW Tau AB, and the 3 σ upper limit for a unresolved source suggests a dust mass ≲0.07 M ⊕ . It is known that close binaries can shorten disk lifetimes (e.g., Cieza et al. 2009 ), so the disk around FW Tau AB may be depleted already. For the companion FW Tau C, its dust disk is compact and likely unresolved; in contrast, the gas disk is more extended and clearly shows a Keplerian rotation. Figure 2 is the position–velocity diagram constructed along the major axis of the gas disk. On top of the PV diagram we also plot the Keplerian rotation curves for 10 M Jup and 0.1 M objects. It is clear that the velocity profile of the FW Tau C disk is incompatible with the 10 M Jup rotation curve, but more consistent with that of a 0.1 M star. Our disk modeling (see Figure 3 and Table 2 ) also indicates that FW Tau C’s mass is ∼0.1 M , about 10 times higher than the 10 M Jup suggested by Kraus et al. ( 2014 ). Hence, FW Tau C is not a planetary-mass object, but is rather a low-mass star with an inclined disk, as suggested by Bowler et al. ( 2014 ). The high mass of FW Tau C is also in agreement with some features in its spectrum that closely resemble T Tauri stars (Bowler et al. 2014 ). Our observations therefore suggest that FW Tau is a young triple system in which each of the stars has a mass of ∼0.1 M . Figure 2. Position–velocity diagram constructed along the major axis of the FW Tau C disk. Keplerian rotation curves for a 0.01 M (∼10 M Jup ) planet and a 0.1 M star are also plotted. FW Tau C is more consistent with a 0.1 M star. Figure 3. We show our CO (2–1) channel maps (top) as well as the best-fit Keplerian disk model synthetic channel maps (center) and the residual channel maps (bottom). The model channel maps are made by sampling the model at the same baselines as the data and Fourier transforming to produce images. The residual channel maps are produced by subtracting the model visibilities from the visibility data and Fourier transforming to produce maps. The contours start at 3 σ = 13.5 mJy beam −1 and are at intervals of 3 σ . Table 2 FW Tau C Disk Properties Searches for accretion disks in millimeter around directly imaged planet-mass companions have not yielded positive results (e.g., Isella et al. 2014 ; Bowler et al. 2015 ; MacGregor et al. 2017 ; Ricci et al. 2017 ; Wolff et al. 2017 ; Wu et al. 2017 ), although the disk around the free-floating planet OTS 44 has recently been imaged (Bayo et al. 2017 ). Current flux upper limits imply that these wide companion disks might have 0.1 M ⊕ of dust, which in turn could imply a
Paper: 2017MNRAS.468.1594A (acknowledgments)
Acknowledgments The authors are grateful to Ray Sharples for the use of KMOS GTO time for some observations and for useful discussions regarding instrumentation; our thanks also to Renske Smit and Helen Johnson for obtaining additional observations of NGC1407 and NGC3379 for this project. We wish to thank Charlie Conroy for providing up-to-date versions of the CvD models and the anonymous referee for helpful comments that improved the paper. PA was supported by an STFC studentship (ST/K501979/1) and RJS by the STFC Durham Astronomy Consolidated Grant 20142017 (ST/L00075X/1). This work is based on observations obtained at the Very Large Telescope of the European Southern Observatory (programme IDs: 093.B-0693 and 094.B-0061). The data used here are available through the ESO science archive (programme IDs: 093.B-0693(A), 094.B-0061(C)).
Paper: 2019ApJ...874...65D (acknowledgments)
We gratefully acknowledge the support of the STScI team responsible for ensuring timely and accurate implementation of our M87 programs. Support for programs #10543 and #14618 was provided by NASA through a grant from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. PyRAF is a product of the Space Telescope Science Institute, which is operated by AURA for NASA.
Paper: 2016MNRAS.458.3118B (fulltext)
In two cases, additional variable stars are found (rms significantly larger than for the remaining objects), they are excluded from the sample of comparison stars. A small number of epochs (not more than one per field) produces outliers in the master light curve probably due to changing in atmospheric conditions not recognized during the run; these images were excluded from the analysis. The master light curve was re-calculated using the remaining set of comparison stars and images and again subtracted from the raw light curves. Fig. 2 presents the final differential light curves and rms plots (rms versus average magnitude) for each of our seven objects. In these plots, our targets are labelled. The typical photometric error (defined as the floor of the rms of the comparison stars) is ∼ 1–3 per cent, significantly smaller than the amplitudes in the variable targets. Some of the target light curves exhibit a distinctive morphology – V2737 shows a continuous downward trend, V2739 a periodic behaviour, and V2227 multiple short-period eclipses. The remaining light curves appear at least partially irregular. Although the sampling frequency of the current data set is less dense than in previous studies, the overall light-curve morphology and variability time-scales are comparable to the previously published ones for all objects. Figure 2. Observed light curves and rms plots for all seven objects. Left column contains light curves, right column contains corresponding rms plot for each target's field. The light curves for each object span four nights and shows relative magnitudes from differential photometry. We observe a spread of periodic, quasi-periodic and irregular behaviours. rms plots indicate variable objects in each target's field. Figure 2 – continued Figure 2 – continued We calculate amplitudes for the light curves by taking the min–max difference for each star. To ensure no extreme outliers are included in this calculation, we check for any 3σ outliers but find none. Table 3 lists the derived amplitudes and compares them to previously published values ( SE04 ; SE05 ). The amplitudes in SE04 and SE05 refer to the peak-to-peak values in the binned light curves. Table 3. Light-curve amplitudes ( A ) derived in this work (see Section 6.1 ) and previously published values (see text for details). a SE04 ; table 4, I band. b SE05 ; table 1, I band. 4.2 Spectroscopy Standard spectroscopic data reduction was performed with iraf routines from the onedspec and twodspec packages. Background was fit and subtracted during spectrum extraction. Wavelength calibration is done using an HeAr arclamp. The A-type star HD292956 was used for relative flux calibration. For each epoch, three spectra (see Section 3 ) were co-added using a median filter. Fig. 3 presents the complete sequences of calibrated spectra for our stars over the four nights of observations. Every spectrum is normalized to the central wavelength of the I -band filter (7680 Å). In the final
Paper: 2020MNRAS.497.4602F (fulltext)
can also be seen in more detail in Fig. 9 , where we show the running median spot properties and the typical uncertainties in spot coverage versus spot temperature plots. Phase 1: a general decline of amplitudes in V and R c , while remains almost constant. During the first half, the spot temperature declines from about 5500 to 5300 K at constant spot coverage (0.3), followed by a shrinking of the spot coverage to 0.2 at constant temperature. The initial decrease of temperature occurs at a rate of about −2.2 K/d. Overall, this is a phase where spot activity declines. Phase 2: there are only minor variations in the amplitudes, with a maximum in roughly in the middle of the phase. The spot temperatures are constant (5450 K) within the uncertainties, with a weak trend of slightly increasing values at a rate of about +0.2 K/d. The spot coverage (0.2) is also constant within the uncertainty. Phase 3: there is a clear trend of decreasing amplitudes in all filters, by at least a factor of 2. As for the previous phase the spot temperatures do remain constant within the uncertainties at about 5600 K. The spot coverage decreases from 0.35 to 0.2, with a small increase towards the end of this phase. This phase contains the high cadence data (indicated as grey in Fig. 8 ). Due to the large amount of data the amplitude uncertainties are much lower, and hence also the spot temperature and coverage uncertainties. The analysis of this data with a much smaller ±15 d window is shown in the bottom left panel of Fig. 6 . The results are consistent with the results obtained for the entire light curve. They show the spot temperature decreasing from 5650 to 5500 K. At the same time, the spot coverage drops from 0.35 to 0.2. Overall, the characteristics of phase 3 are similar to phase 1, but at slightly higher spot temperatures. Phase 4: the amplitudes remain constant for a large part, but start to slowly increase towards the end. Accordingly the spot temperature (5400 K) and coverage (0.15) remain constant throughout, and only increase towards the end. The temperature increases at a rate of +2.9 K/d towards the end. We note that the duration of all phases is significantly longer than our time window of ±60 d, and they are hence unlikely to be an artefact. One possibility to interpret these phases is as lifetime of large spots or groups of spots. If that is the correct interpretation, the spot lifetime in the four phases would be >180, 230, 360, and >230 d, overall, or between half a year and one year. These values are consistent with spot lifetimes determined in the literature for active stars. For example, Giles, Collier Cameron Haywood ( 2017 ) used Kepler data to estimate spot decay time-scales between 20 and 300 d for K-stars. A fast rotating single star like V1598 Cyg would be expected to have a shorter spot lifetime, but a potential close companion and tidal coupling may extend the lifetime. As described in Section 4.1 , we find a long-term change in the
Paper: 2019A&A...625A.148D (acknowledgments)
Based on observations carried out with the IRAM Plateau de Bure Interferometer. IRAM is supported by INSU/CNRS (France), MPG (Germany), and IGN (Spain). We acknowledge the IRAM Plateau de Bure team for the observations. We thank Sabine König for her support in the data reduction. We warmly thank Eric Rosolowsky, our referee, for his very helpful and constructive comments. We are grateful to Pauline Barmby for providing us with reduced Spitzer data in the nucleus region. ALM has benefited from the support of Action Fédératrice Structuration de l’Univers et Cosmologie from Paris Observatory. This research has made use of the NASA/IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. This research made use of Montage, funded by the National Aeronautics and Space Administration’s Earth Science Technology Office, Computational Technnologies Project, under Cooperative Agreement Number NCC5-626 between NASA and the California Institute of Technology. The code is maintained by the NASA/IPAC Infrared Science Archive. This research makes use of GALEX, SDSS, 2MASS, Spitzer , and WISE archive data of Andromeda.
Paper: 2018A&A...614A..42M (fulltext)
triggering of the radio nuclear activity, and detect in-falling molecular clouds that may accrete onto the SMBH. The ALMA observations of the 12 CO (2–1) line reveal that multiple clouds are distributed in a complex warped disk. Between 7.5 kpc and 2.5 kpc from the central radio AGN, the disk is oriented as the dust lane of the galaxy in the north-south direction ( PA ~ 180°). At inner radii (2.5 kpc – 700 pc), the disk changes slightly in orientation ( PA ~ 150°) and molecular clouds are found in proximity of dust ( Fig. 2 ) and of warm H 2 gas (see the left panel of Fig. 9 ). In the very central regions ( r ~ 700 pc), the disk changes abruptly its orientation and the molecular clouds form an edge-on circumnuclear disk oriented approximately east-west (PA ~ 72°). The carbon monoxide is a tracer of the cold phase of the molecular hydrogen. Depending on the conversion factor assumed, the total mass of cold molecular hydrogen is M (H 2 ) 3.2 × 10 8 –1.9 × 10 9 M ⊙ (see Table 3 ), while the circumnuclear disk has total mass of cold molecular hydrogen between 5.3 × 10 7 M ⊙ and 3.1 × 10 8 M ⊙ . The circumnuclear disk of molecular gas of PKS B1718–649 may form the reservoir of fuel for the radio emission. Although the rotation of the gas follows what is expected from the mass of the central SMBH and the distribution of the stellar body, the velocity dispersion of the gas is higher in the centre of the disk ( r 150 pc, 0.45′′) than at larger radii ( Fig. 7 ). This suggests that close to the radio source the gas may have non-circular orbits. The circumnuclear disk is also traced by the H 2 1-0 S(1) line emission ( Maccagni et al. 2016 ). The ratio between the column density of the H 2 gas traced by the H 2 1-0 S(1) line and by the 12 CO (2–1) line increases by a factor ten in the innermost 75 pc of the circumnuclear disk with respect to its outer regions (~ 110 and ~ 16, respectively). This suggests that, in the centre, the physical conditions of the molecular gas are different. It is possible that the radiative field of the AGN is photo-dissociating the cold component of the molecular gas, traced by the 12 CO (2–1) line, before the warmer component, traced by the H 2 1-0 S(1) line. This would indicate that the newly born AGN is actively changing the physical conditions of the surrounding ISM. Against the 230 GHz emission of the radio AGN, we detect 12 CO (2–1) gas in absorption (see Sect. 3.3 ). The absorption line has a width ( FWHM = 54 km s −1 ) much larger than the typical velocity dispersion of molecular clouds. Hence, either the absorption is tracing multiple clouds, or a single cloud with non-circular orbits, or a combination of both. The absorption line is detected at red-shifted velocities with respect to the systemic velocity (Δ v = +365 ± 20 km s −1 ). Its kinematics deviate from the rotation of the circumnuclear disk and the gas may be falling towards the radio source. At the same red-shifted velocities we also detect in emission a