adsabs/AstroCLIMB huggingface dataset:
Paper: Nondetection of Helium in the Hot Jupiter WASP-48b
Authors: Katherine A. Bennett, Seth Redfield, Antonija Oklopčić, Ilaria Carleo, Joe P. Ninan, Michael Endl
Figure caption: Normalized helium absorption vs. log ﹩{R}_{\mathrm{HK}}^{{\prime} }﹩, with spectral type designated by color. Vertical error bars are depicted for all data points, though some are quite small and are not visible. Error bars on log ﹩{R}_{\mathrm{HK}}^{{\prime} }﹩ are included where available. Nondetections (upper limits on helium absorption) are depicted with vertical downward arrows. There appears to be a slight positive correlation between ﹩{\delta }_{{R}_{P}}/{H}_{\mathrm{eq}}﹩ and log ﹩{R}_{\mathrm{HK}}^{{\prime} }﹩, hinting at a relationship between helium absorption and stellar activity. WASP-48b is in the lower left corner at a ﹩{\delta }_{{R}_{P}}/{H}_{\mathrm{eq}}﹩ value of ∼16.3 and log ﹩{R}_{\mathrm{HK}}^{{\prime} }﹩ value of −5.135.
Paper: Spectroscopic Follow-up of Gaia Exoplanet Candidates: Impostor Binary Stars Invade the Gaia DR3 Astrometric Exoplanet Candidates
Authors: Marcus L. Marcussen, Simon H. Albrecht
Figure caption: NOT targets. Each row corresponds to one of the three NOT fast-track targets. These systems are, from top to bottom: Gaia DR3 052469973468984192, Gaia DR3 916454200349735680, and Gaia DR3 5122670101678217728. Left column: the most likely orbits from the Gaia astrometric two-body solution are shown alongside the orbital parameters and their uncertainties. Middle column: RV curve of the primary component as inferred from the Gaia orbital solution. The left y-axis scaling assumes the companion to have a mass equal to the primary. The right y-axis scaling assumes the companion to be dark. The two NOT observations are shown as solid and dotted lines. The gray shaded areas indicate the 1σ uncertainties in orbital phase of these observations. The rectangles are darker where the uncertainties overlap. The large uncertainties in orbital phase timing are caused by the long baseline between the Gaia and NOT observations. Right column: CCFs obtained from the FIES spectra, using line styles corresponding to those from the middle column.
Paper: Spectral Energy Distributions for 258 Local Volume Galaxies
Authors: Daniel A. Dale, Médéric Boquien, Jordan A. Turner, Daniela Calzetti, Robert C. Kennicutt, Janice C. Lee
Figure caption: The distributions of Bayesian-based fit values for several key parameters are plotted in blue. The red histograms overlaid for the dust parameters γ, q PAH, and ﹩{U}_{\min }﹩ show the values from the spectral analysis of the SINGS sample of nearby galaxies in Draine et al. (2007). The magenta histogram for A Hα  shows the distribution for the 11HUGS survey that has large overlap with LVL (Lee et al. 2009b). Table 3 provides the median values and the 25th–75th percentile ranges for these parameters.
Paper: Spectral Energy Distributions for 258 Local Volume Galaxies
Authors: Daniel A. Dale, Médéric Boquien, Jordan A. Turner, Daniela Calzetti, Robert C. Kennicutt, Janice C. Lee
Figure caption: Top: the fraction of stellar mass in the modeled late burst of star formation compared to the total stellar mass in the galaxy, as a function of total stellar mass. Bottom: the ratio of the mass in the young stellar population (<10 Myr) to the total mass of the stellar population. Filled circles correspond to LVL galaxies; those in green indicate blue compact dwarf galaxies in the LVL sample. The green open circles are from the Janowiecki et al. (2017) survey of blue compact dwarf galaxies and the red open squares are from the Bergvall et al. (2016) study of SDSS starbursting galaxies. Note that the upper panel is restricted to late bursts and not late quenching episodes (i.e., r SFR > 1).
Paper: Spectral Energy Distributions for 258 Local Volume Galaxies
Authors: Daniel A. Dale, Médéric Boquien, Jordan A. Turner, Daniela Calzetti, Robert C. Kennicutt, Janice C. Lee
Figure caption: The (Bayesian-based) star-formation rate from the SED fit averaged over 100 Myr compared to the star-formation rate derived from hybrid indicators of ultraviolet + total infrared (Hao et al. 2011). The left-hand panel is plotted as a function of the 3.6 μm—based stellar mass and the right-hand panel is shown as a function of hybrid-based star-formation rates. The horizontal dotted line indicates unity or perfect agreement between the star-formation rates.
Paper: Spectral Energy Distributions for 258 Local Volume Galaxies
Authors: Daniel A. Dale, Médéric Boquien, Jordan A. Turner, Daniela Calzetti, Robert C. Kennicutt, Janice C. Lee
Figure caption: The fractional mass abundance of PAHS vs. minimum interstellar radiation field intensity (left) and the fractional PAH mass abundance from the SED fitting work of Draine et al. (2007) for the subset of the LVL sample that overlaps with the SINGS sample (right). The median values for several bins are provided in the left-hand panel, the dotted line in the right-hand panel indicates the maximum possible value for the Draine et al. (2007) SED fits, and the dashed line in the right-hand panel indicates the 1-to-1 relation.
Paper: Spectroscopic Follow-up of Gaia Exoplanet Candidates: Impostor Binary Stars Invade the Gaia DR3 Astrometric Exoplanet Candidates
Authors: Marcus L. Marcussen, Simon H. Albrecht
Figure caption: Astrometric exoplanet candidate systems in an H-R diagram. All Gaia DR3 systems where the astrometric two-body motion of the source is small enough to indicate a companion with a mass <20M Jup, assuming the companion does not contribute any flux, are shown here as circles. The released Gaia astrometric exoplanet candidates are shown using closed circles while open circles indicate systems where a lower bound on the flux of the companion is nonzero as per the binary_masses table. Large circles indicate systems where the nature of the companion is confirmed, either here or in the literature; see also Table 1. The systems are colored according to the current knowledge of the nature of the companion, where brown indicates a planetary companion, orange indicates a stellar companion, and cyan indicates that it is unknown. The two white dots indicate two white dwarf hosts where the companion nature is unknown. The comparison stars shown in the background are the 1,000,000 nearest-to-us Gaia DR3 stars whose estimated parallax is more than 10 times its uncertainty. The density of these stars is represented as a purple-to-yellow heatmap, where purple (yellow) indicates a low (high) density. Since most of the planet candidates are too faint to have Gaia DR3 reddening and extinction parameters, candidates and comparison stars alike are shown using their uncorrected colors and magnitudes. The black dashed line represents the best fit to the main sequence. The six systems discussed in detail in Section 4 are indicated by their name.
Paper: Spectroscopic Follow-up of Gaia Exoplanet Candidates: Impostor Binary Stars Invade the Gaia DR3 Astrometric Exoplanet Candidates
Authors: Marcus L. Marcussen, Simon H. Albrecht
Figure caption: Photocenter. Not-to-scale illustration of three different geometries of face-on, circular orbits that the same astrometric signal may originate from. Left panel: an exoplanet orbiting a star. The semimajor axis of the exoplanet orbit, a 2, is much larger than the orbit of the star, a 1, due to the large difference in mass. The flux contribution from the planet is negligible, and thus the photocenter closely follows the star. Middle panel: twin stars (i.e., almost equal flux and mass stars) orbit their barycenter (black dot). The primary component, the yellow star, is slightly more massive than the secondary component, the red star, leading to a smaller semimajor axis a 1 than a 2. The ratio of the masses is slightly higher than the ratio of the received fluxes, M 2/M 1 > F 2/F 1, resulting in a small photocenter semimajor axis a 0 that traces out (a diminished version of) the orbit of the primary component. Right panel: two stars with significantly different masses and fluxes in a hierarchical triple system, where the center-of-mass of the red and yellow star orbits around a larger and brighter third component that is very far away and shown in white. The flux of the third component diminishes the size of the photocenter semimajor axis but does not contribute to the astrometric signal due to its very long period. The photocenter semimajor axis would be larger and would more closely follow the yellow star if not for the third component.
Paper: Spectroscopic Follow-up of Gaia Exoplanet Candidates: Impostor Binary Stars Invade the Gaia DR3 Astrometric Exoplanet Candidates
Authors: Marcus L. Marcussen, Simon H. Albrecht
Figure caption: HD 68638. Top left: the geometry of the Gaia two-body orbital solution for the photocenter is shown with north pointing up and east pointing left. The Gaia parameters are printed as well. The barycenter is shown as a black dot, and the star is shown at periastron as an orange dot with the direction of motion indicated by the black arrow. The line through the barycenter represents the orbital angular momentum vector. The positive direction of the angular momentum is marked with an arrow that is occluded by the orbital disk in this particular view. The light gray part of the orbit with a solid orange outline indicates the near side of the sky-plane, and the dark gray part with a dashed outline indicates the far side of the sky-plane. The blue circles are the location on the orbit of the RV measurements. Bottom left: CCFs from all 27 publicly available ELODIE spectra are shown and colored according to their FWHM. Top right: the solid black line indicates the expected radial-velocity curve of the primary component based on the Gaia-derived orbital and stellar parameters assuming two equal-mass companions. The dashed line represents an orbit with a higher eccentricity of 0.6. Middle right: the FWHM of the CCFs is shown as blue dots as a function of orbital phase. The solid black line is the FWHM as a function of orbital phase inferred from the Gaia orbit. The dashed black line represents the FWHM of a more eccentric orbit that is a better fit to the observed FWHM. Bottom right: CCFs plotted with a horizontal offset, so their midpoints are located under their respective FWHM data points of the middle right panel.
Paper: Spectroscopic Follow-up of Gaia Exoplanet Candidates: Impostor Binary Stars Invade the Gaia DR3 Astrometric Exoplanet Candidates
Authors: Marcus L. Marcussen, Simon H. Albrecht
Figure caption: Exoplanet candidate parameters. Top panel: companion masses of the same sample as shown in Figure 1 are calculated under the assumption that their flux contribution is negligible compared to their host and assuming Gaussian uncertainties in photocenter semimajor axis a0, orbital period P, and primary mass M 1. Middle panel: the G magnitudes of the sample. Bottom panel: the size of the photocenter semimajor axis relative to its uncertainty, which can be taken as a proxy for the robustness of the Keplarian orbit detection.
Paper: Spectroscopic Follow-up of Gaia Exoplanet Candidates: Impostor Binary Stars Invade the Gaia DR3 Astrometric Exoplanet Candidates
Authors: Marcus L. Marcussen, Simon H. Albrecht
Figure caption: Astrometry and RV simulations. Left: the mass of a companion as determined by astrometry (dashed line) and RVs (solid lines) relative to its actual mass as a function of the flux ratio divided by the mass ratio, ϵ/q = (F 2/F 1)/(M 2/M 1). The line colors indicate the ratio of the RV semiamplitude to the absorption line widths, ﹩v\sin i﹩. Right: the ratio of the RV-determined mass to the astrometry-determined mass.
Paper: Spectroscopic Follow-up of Gaia Exoplanet Candidates: Impostor Binary Stars Invade the Gaia DR3 Astrometric Exoplanet Candidates
Authors: Marcus L. Marcussen, Simon H. Albrecht
Figure caption: HIP 66074. Shown here are the Gaia geometry and orbital parameters of HIP 66074 as well as RV measurements in blue derived from HIRES spectra. The Gaia astrometry and RV data are in excellent agreement except for a large factor 15 disagreement in the RV amplitudes between the data sets. In the bottom left, we show CCFs that we obtained from the HIRES spectra. In contrast to Figure 5, the CCFs from the different observations are so similar that they look like one, ruling a stellar companion.
Paper: Spectroscopic Follow-up of Gaia Exoplanet Candidates: Impostor Binary Stars Invade the Gaia DR3 Astrometric Exoplanet Candidates
Authors: Marcus L. Marcussen, Simon H. Albrecht
Figure caption: HD 40503. Shown here are the geometry and orbital parameters of HD 40503 as well as RV measurements in red and blue from the HARPS and CORALIE spectrographs respectively. The black dashed line represents the best fit to the RV data alone.
Paper: Single-band VLBI Absolute Astrometry
Authors: Leonid Petrov
Figure caption: The rms of the differences in right ascension (upper row) and declination (lower row) when single-band group delays are used with respect to the positions from the dual-band reference solution. The a priori ionospheric path delay using GNSS TEC with the modified mapping function (ΔH = 56.7 km, α = 0.9782, and k = 0.85) was used by adjusting the TEC biases from the dual-band observations (right) and without adjusting the biases but with applying the debias correction from Equation (13), shown in the left column. The upper blue band shows the differences for S band, the next red line shows the differences for C band, the next green line shows the differences for X band, and the bottom purple line shows the differences for K band.
Paper: Single-band VLBI Absolute Astrometry
Authors: Leonid Petrov
Figure caption: The rms of the mean residual ionospheric contribution at VLBA baselines at 8 GHz for three cases: (1) no a priori ionospheric contribution is applied (upper red curve); (2) the a priori ionospheric path delay computed using GNSS TEC maps is applied (middle blue curve); and (3) the a priori ionospheric path delay computed using GNSS TEC maps are applied and the TEC biases are adjusted (lower green curve). 1 TECU causes a group delay of 21 ps at 8 GHz.
Paper: Single-band VLBI Absolute Astrometry
Authors: Leonid Petrov
Figure caption: Differences in declination from the X-band solution XIT using data from the VLBA network with the data reduction for the ionosphere using the ionospheric path delay from the GNSS TEC maps applied with respect to the dual-band reference solution. The thick blue line shows the differences smoothed with a Gaussian filter.
Paper: Single-band VLBI Absolute Astrometry
Authors: Leonid Petrov
Figure caption: An increase in source position errors in right ascension (left) and declination (right) due to the residual ionospheric contribution in K band (22 GHz) after applying data reduction based on the GNSS TEC maps with the modified mapping function (ΔH = 56.7 km, α = 0.9782, and k = 0.85).
Paper: Single-band VLBI Absolute Astrometry
Authors: Leonid Petrov
Figure caption: Smoothed declination biases from four solutions with respect to the dual-band reference solution using X-band observables and the ionospheric mapping function with parameters ΔH = 150 km and α = 0.9782 with different scaling factors k = 0.7, 0.8, 0.9, and 1.0.
Paper: Single-band VLBI Absolute Astrometry
Authors: Leonid Petrov
Figure caption: Smoothed declination biases from four solutions with respect to the dual-band reference solution using X-band observables and the ionospheric mapping function with parameters ΔH = 0, α = 1.0 with different scaling factors k = 0.7, 0.8, 0.9, and 1.0.
Paper: Single-band VLBI Absolute Astrometry
Authors: Leonid Petrov
Figure caption: Dependence of the VLBI ionospheric group delay at 8 GHz against the ionospheric group delay from GNSS. The blue straight line is the least-squares fit of this dependence.
Paper: Single-band VLBI Absolute Astrometry
Authors: Leonid Petrov
Figure caption: The declination with the ionospheric path delay using the mapping function with parameters ΔH = 56.7 km, α = 0.9782, and k = 0.85 with (right) and without (left) applying the empirical debias correction.
Paper: Single-band VLBI Absolute Astrometry
Authors: Leonid Petrov
Figure caption: Smoothed declination biases from four solution with respect to the dual-band reference solution using X-band observables and the ionospheric mapping function with parameters ΔH = 56.7 km and α = 0.9782 with different scaling factors k = 0.7, 0.8, 0.9, and 1.0.
Paper: Single-band VLBI Absolute Astrometry
Authors: Leonid Petrov
Figure caption: The distribution of the normalized differences of the ionospheric path delay computed from the GNSS TEC maps against the VLBI ionospheric path delay with clock biases subtracted (green dots). The normal distribution with σ = 1 (solid blue line) is shown as a reference.
Paper: Single-band VLBI Absolute Astrometry
Authors: Leonid Petrov
Figure caption: The distribution of normalized differences in declination from the trial VLBI solutions with respect to the reference dual-band solution. Left: solution with GNSS TEC maps + adjustments of the TEC biases using dual-band VLBI group delays (XIA). Right: solution with added Gaussian noise with σ equal to the errors in the path delay, which correspond to the TEC bias adjustments uncertainties (XIN). The thick blue lines show the fitted Gaussian distributions with the first and second moments listed in Table 1.
Paper: Single-band VLBI Absolute Astrometry
Authors: Leonid Petrov
Figure caption: Empirical distribution of the normalized differences of the ionosphere path delay computed from the GNSS TEC maps adjusted for clock and TEC biases (green dots). The normal distribution with σ = 1 is shown as a reference (solid blue line).
Paper: Single-band VLBI Absolute Astrometry
Authors: Leonid Petrov
Figure caption: Dependence of the rms of the residual ionospheric path delay derived from GNSS TEC maps σ gr on the rms of the total ionospheric path delay from these maps σ gt. No adjustment to the TEC has been applied. The path delay is computed for the reference frequency 8 GHz. The blue smooth line shows the regression model in a form of a B-spline that fits the data.
Paper: Single-band VLBI Absolute Astrometry
Authors: Leonid Petrov
Figure caption: The distribution of second-moment estimates of the normalized differences of the ionospheric path delays derived from VLBI dual-band observations and GNSS TEC maps among individual observing sessions. The narrow green curve shows the statistics of the normalized residuals with TEC biases adjusted and the wide blue curve show the statistics of the normalized residuals without TEC adjustment.
Paper: Single-band VLBI Absolute Astrometry
Authors: Leonid Petrov
Figure caption: Adjustment to the ionosphere path delay bias at 8.4 GHz with respect to the path delay derived from GNSS TEC maps at the MK-VLBA station from processing dual-band observations on 2015 April 22. The bias estimate corresponds to the averaged positions of the ionosphere piercing points within a given time interval. 1 TECU causes a group delay of 21 ps at 8.4 GHz.
Paper: Dynamical Evolution of Closely Packed Multiple Planetary Systems Subject to Atmospheric Mass Loss
Authors: Su Wang, 素 王, D. N. C. Lin, 潮 林
Figure caption: The evolution of the system in Case 1, where A = 0.05, and τ  m  = 106 yr. Panel (a) shows the evolution of the innermost planet. The black line represents the evolution of the semimajor axis, and the red line represents the evolution of mass. In this case, the innermost planet will lose about 0.14 Earth-mass, which is about 5% of the initial mass. Panel (b) shows the evolution of semimajor axis, pericenter, and apocenter. The black lines show the evolution of the semimajor axis, and the gray lines show the evolution of the pericenter and apocenter. At the end of the simulation, the eccentricities of the planets will be excited. However, the system is still stable due to the relative angles between them. Panel (c) shows the evolution of the resonant angles θ  ji  of two adjacent planets  Pi  and P  j , 1 ≤ i ≤ 8, j = i + 1. ﹩{\theta }_{{ji}}=(p+q)\lambda ^{\prime} -p\lambda -q\varpi ^{\prime} ﹩, where the unprimed and primed parameters refer to the inner and outer planets, respectively, and two planets are in or near (p+q):p resonance. λ and ϖ are the mean longitude and longitude of pericenter. Panel (d) shows the evolution of eccentricities of planets in the system e  pi , where pi represents the serial number of the planet from inner to outer, 1 ≤ pi ≤ 9, and the eccentricities of different planets are labeled in different colors, as shown in the legend in panel (d). The evolution with A = 0.1 and τ  m  = 106 yr in Case 2 is similar to that in Case 1.
Paper: Dynamical Evolution of Closely Packed Multiple Planetary Systems Subject to Atmospheric Mass Loss
Authors: Su Wang, 素 王, D. N. C. Lin, 潮 林
Figure caption: The distribution of crossing time vs. the separation change of planet pairs. The original system that is obtained from G2 can be stable for more than 107 yr, as shown in Figure 5 . In order to find out the stable region for such multiple planetary system, we introduce an artificial factor f  k  acting on the original stable system to change the separation between adjacent planets. We fix the semimajor axis of the middle planet P 5 and change the separation between adjacent planets by f  k , 0.8 ≤ f  k  ≤ 1.25. f  k  < 1 refers to a global contraction comparing to the original system configuration, and f  k  > 1 represents an expansion to the original stable system.
Paper: Dynamical Evolution of Closely Packed Multiple Planetary Systems Subject to Atmospheric Mass Loss
Authors: Su Wang, 素 王, D. N. C. Lin, 潮 林
Figure caption: The evolution of the system in Case 3, where A = 0.05 and τ  m  = 104 yr. Panel (a) shows the evolution of the innermost planet. In this case, the innermost planet will lose about 0.14 Earth-mass, which is about 5% of the initial mass. Panel (b) shows the evolution of a, a(1-e), and a(1+e) of all of the planets. At the end of the simulation, the results are similar to those in Case 3. Panel (c) shows the evolution of the resonant angles. The resonant angles tend to librate in a larger range than in Case 1. Panel (d) shows the evolution of eccentricities of planets in the system. The different colors represent the same as those in Figure 9.
Paper: Dynamical Evolution of Closely Packed Multiple Planetary Systems Subject to Atmospheric Mass Loss
Authors: Su Wang, 素 王, D. N. C. Lin, 潮 林
Figure caption: The evolution of the system with five planets, where A = 0.25 and τ  m  = 104 yr. The system becomes unstable at about 0.06 Myr. Panel (a) shows the evolution of the innermost planet. After the system become unstable, the inner three planets collide into one larger planet, which is about 9.33 M ⊕. And then planets 4 and 5 collide into another single planet with 6.67 M ⊕. Panel (b) shows the evolution of a, a(1-e), and a(1+e) of all of the planets. Panel (c) shows the evolution of the resonant angles. Panel (d) shows the evolution of eccentricities of planets in the system. The different colors represent the same as those in Figure 9.
Paper: Dynamical Evolution of Closely Packed Multiple Planetary Systems Subject to Atmospheric Mass Loss
Authors: Su Wang, 素 王, D. N. C. Lin, 潮 林
Figure caption: The evolution of the system with five planets, where A = 0.15 and τ  m  = 104 yr. Panel (a) shows the evolution of the innermost planet. In this case, the innermost planet will lose about 0.5 Earth-mass, which is about 15% of the initial mass. Panel (b) shows the evolution of a, a(1-e), and a(1+e) of all of the planets. At the last stage of the simulation, the planet pairs in the system tend to escape from the MMRs. Panel (c) shows the evolution of the resonant angles. The resonant angles tend to librate in a large range at the end of simulation. Panel (d) shows the evolution of eccentricities of planets in the system. The different colors represent the same as those in Figure 9.