adsabs/AstroCLIMB huggingface dataset:
Paper: A High-resolution Spectroscopic Survey of Directly Imaged Companion Hosts. I. Determination of Diagnostic Stellar Abundances for Planet Formation and Composition
Authors: Aneesh Baburaj, Quinn M. Konopacky, Christopher A. Theissen, Sarah Peacock, Lori Huseby, Benjamin J. Fulton, Roman Gerasimov, Travis S. Barman, Kielan K. W. Hoch
Figure caption: Best-fit PHOENIX-C/O model of the APF spectrum for the target HD 206893 (cyan), shown for order #76 containing an O I triplet feature at 6155–6158 Å. The zoomed-in subplot on the right shows a ∼10 Å region around the blended feature. The best-fit PHOENIX-C/O model has T eff = 6617 K, ﹩\mathrm{log}g﹩ = 4.27, ﹩\mathrm{log}{\epsilon }_{{\rm{C}}}﹩ = 8.57, and ﹩\mathrm{log}{\epsilon }_{{\rm{O}}}﹩ = 8.66 (magenta). The residuals between the data and the model are plotted in black, and other noise limits are shown in gray.
Paper: A High-resolution Spectroscopic Survey of Directly Imaged Companion Hosts. I. Determination of Diagnostic Stellar Abundances for Planet Formation and Composition
Authors: Aneesh Baburaj, Quinn M. Konopacky, Christopher A. Theissen, Sarah Peacock, Lori Huseby, Benjamin J. Fulton, Roman Gerasimov, Travis S. Barman, Kielan K. W. Hoch
Figure caption: Best-fit PHOENIX-C/O model of the APF spectrum for the target HD 984 (cyan), shown for order #76 containing an O I triplet feature at 6155–6158 Å. The zoomed-in subplot on the right shows a ∼10 Å region around the blended feature. The best-fit PHOENIX-C/O model has T eff = 6401 K, ﹩\mathrm{log}g﹩ = 4.42, ﹩\mathrm{log}{\epsilon }_{{\rm{C}}}﹩ = 8.49, and ﹩\mathrm{log}{\epsilon }_{{\rm{O}}}﹩ = 8.69 (magenta). The residuals between the data and the model are plotted in black, and other noise limits are shown in gray.
Paper: A High-resolution Spectroscopic Survey of Directly Imaged Companion Hosts. I. Determination of Diagnostic Stellar Abundances for Planet Formation and Composition
Authors: Aneesh Baburaj, Quinn M. Konopacky, Christopher A. Theissen, Sarah Peacock, Lori Huseby, Benjamin J. Fulton, Roman Gerasimov, Travis S. Barman, Kielan K. W. Hoch
Figure caption: Best-fit PHOENIX-C/O model of the APF spectrum for the target HR 8799 (cyan), shown for order #76 containing an O I triplet feature at 6155–6158 Å. The zoomed-in subplot on the right shows a ∼10 Å region around the blended feature. The best-fit PHOENIX-C/O model has T eff = 7317 K, ﹩\mathrm{log}g﹩ = 4.32, ﹩\mathrm{log}{\epsilon }_{{\rm{C}}}﹩ = 8.56, and ﹩\mathrm{log}{\epsilon }_{{\rm{O}}}﹩ = 8.79 (magenta). The residuals between the data and the model are plotted in black, and other noise limits are shown in gray.
Paper: A High-resolution Spectroscopic Survey of Directly Imaged Companion Hosts. I. Determination of Diagnostic Stellar Abundances for Planet Formation and Composition
Authors: Aneesh Baburaj, Quinn M. Konopacky, Christopher A. Theissen, Sarah Peacock, Lori Huseby, Benjamin J. Fulton, Roman Gerasimov, Travis S. Barman, Kielan K. W. Hoch
Figure caption: Corner plot for one of the retained runs for a multi-order fit of the PHOENIX grid to the spectrum of HD 206893. The marginalized posteriors are shown along the diagonal. The blue lines represent the 50 percentile, and the dotted lines represent the 16 and 84 percentiles. The subsequent covariances between all the parameters are in the corresponding 2D histograms. This run gives a best-fit T eff = 6640 K, ﹩\mathrm{log}g﹩ = 4.28, and [M/H] = 0.04.
Paper: A High-resolution Spectroscopic Survey of Directly Imaged Companion Hosts. I. Determination of Diagnostic Stellar Abundances for Planet Formation and Composition
Authors: Aneesh Baburaj, Quinn M. Konopacky, Christopher A. Theissen, Sarah Peacock, Lori Huseby, Benjamin J. Fulton, Roman Gerasimov, Travis S. Barman, Kielan K. W. Hoch
Figure caption: Corner plot for one of the retained runs for the PHOENIX-C/O grid fit to the spectrum of HD 984. The marginalized posteriors are shown along the diagonal. The blue lines represent the 50 percentile, and the dotted lines represent the 16 and 84 percentiles. The subsequent covariances between all the parameters are in the corresponding 2D histograms. This run gives a best-fit ﹩\mathrm{log}{\epsilon }_{{\rm{C}}}﹩ = 8.44, ﹩\mathrm{log}{\epsilon }_{{\rm{O}}}﹩ = 8.71.
Paper: A High-resolution Spectroscopic Survey of Directly Imaged Companion Hosts. I. Determination of Diagnostic Stellar Abundances for Planet Formation and Composition
Authors: Aneesh Baburaj, Quinn M. Konopacky, Christopher A. Theissen, Sarah Peacock, Lori Huseby, Benjamin J. Fulton, Roman Gerasimov, Travis S. Barman, Kielan K. W. Hoch
Figure caption: Best-fit PHOENIX-C/O model of the APF spectrum for the target HIP 25278 (cyan), shown for order #76 containing an O I triplet feature at 6155–6158 Å. The zoomed-in subplot on the right shows a ∼10 Å region around the blended feature. The best-fit PHOENIX-C/O model has T eff = 6173 K, ﹩\mathrm{log}g﹩ = 4.33, ﹩\mathrm{log}{\epsilon }_{{\rm{C}}}﹩ = 8.52, and ﹩\mathrm{log}{\epsilon }_{{\rm{O}}}﹩ = 8.75 (magenta). The residuals between the data and the model are plotted in black, and other noise limits are shown in gray. While the model disagrees with our spectra, especially around 6130, 6145, and 6160 Å, this disagreement is due to the model Fe I lines not agreeing with our spectra. This is expected, as we do not use the PHOENIX-C/O grid with the target metallicity ([M/H] = −0.09) to fit for the C and O abundances.
Paper: A High-resolution Spectroscopic Survey of Directly Imaged Companion Hosts. I. Determination of Diagnostic Stellar Abundances for Planet Formation and Composition
Authors: Aneesh Baburaj, Quinn M. Konopacky, Christopher A. Theissen, Sarah Peacock, Lori Huseby, Benjamin J. Fulton, Roman Gerasimov, Travis S. Barman, Kielan K. W. Hoch
Figure caption: Best-fit PHOENIX model of the APF spectrum for the target HIP 25278 (cyan), shown for APF order #78. The parameters of the best-fit PHOENIX model are estimated by computing the median for the effective temperature (T eff), surface gravity (﹩\mathrm{log}g﹩), and the metallicity ([M/H]) over seven retained runs after performing multiple multi-order MCMC runs. This model has T eff = 6174 K, ﹩\mathrm{log}g﹩ = 4.33, and [M/H] = −0.09 (magenta). The residuals between the data and the model are plotted in black, and other noise limits are shown in gray.
Paper: A High-resolution Spectroscopic Survey of Directly Imaged Companion Hosts. I. Determination of Diagnostic Stellar Abundances for Planet Formation and Composition
Authors: Aneesh Baburaj, Quinn M. Konopacky, Christopher A. Theissen, Sarah Peacock, Lori Huseby, Benjamin J. Fulton, Roman Gerasimov, Travis S. Barman, Kielan K. W. Hoch
Figure caption: Corner plot for one of the retained runs for PHOENIX-C/O grid fit to the spectrum of HIP 25278. The marginalized posteriors are shown along the diagonal. The blue lines represent the 50 percentile, and the dotted lines represent the 16 and 84 percentiles. The subsequent covariances between all the parameters are in the corresponding 2D histograms. This run gives a best-fit ﹩\mathrm{log}{\epsilon }_{{\rm{C}}}﹩ = 8.57, ﹩\mathrm{log}{\epsilon }_{{\rm{O}}}﹩ = 8.80.
Paper: A High-resolution Spectroscopic Survey of Directly Imaged Companion Hosts. I. Determination of Diagnostic Stellar Abundances for Planet Formation and Composition
Authors: Aneesh Baburaj, Quinn M. Konopacky, Christopher A. Theissen, Sarah Peacock, Lori Huseby, Benjamin J. Fulton, Roman Gerasimov, Travis S. Barman, Kielan K. W. Hoch
Figure caption: Corner plot for one of the retained runs for a multi-order fit of the PHOENIX grid to the spectrum of HD 206893. The marginalized posteriors are shown along the diagonal. The blue lines represent the 50 percentile, and the dotted lines represent the 16 and 84 percentiles. The subsequent covariances between all the parameters are in the corresponding 2D histograms. This run gives a best-fit T eff = 6153 K, ﹩\mathrm{log}g﹩ = 4.33, and [M/H] = −0.11.
Paper: Lessons from the High-resolution Spectroscopy of AW UMa and ϵ CrA: Is the Lucy Model Valid?
Authors: Slavek M. Rucinski
Figure caption: The thin vertical bars show measurements of the edge-to-edge extent of the secondary component RV profiles for ϵ CrA, expressed relative to the velocity of the secondary component's mass center. These repeated and difficult estimates were derived from images formed from individual BF profiles. The broken wavy lines give the expected range of velocities for the inner critical Roche lobe of the mass ratio q = 0.130. Observations from individual nights are color-coded as per the legend in the figure, where the integer part of the phase gives the number of elapsed orbits from the assumed initial epoch (see P2). Thick black lines at the upper and lower figure frames indicate the phase ranges of the estimated total and partial eclipses. In the left part of the figure, the phase drift of the dominant ESP is represented by colored symbols according to the figure legend. The inclined lines alongside these symbols represent the predicted phase drifts for three sub-observer longitudes, expressed in orbital phases: 0.6, 0.65, and 0.7.
Paper: Lessons from the High-resolution Spectroscopy of AW UMa and ϵ CrA: Is the Lucy Model Valid?
Authors: Slavek M. Rucinski
Figure caption: The BF profiles for ϵ CrA at both orbital quadratures, i.e., phases ϕ = 0.25 (red) and ϕ = 0.75 (blue), in units per an RV sampling interval of 8.5 km s−1. The profile ϕ = 0.75 is inverted in the velocity scale with the common origin at the RV of the binary mass center at V0 = −62.54 km s−1. The Lucy-model profile for the orbital phase ϕ = 0.25, the assumed mass ratio q = 0.13, and the fill-out parameter f = 0.25 is shown by a thin black line. Vertical dashes mark the mass-center RVs while a thin, broken line shows the rotational profile fit to the primary component. The red and blue bars over the secondary component show the widths of the secondary profile observed at both orbital quadratures. An ESP affecting the secondary profile at ϕ = 0.75 is indicated by an arrow at 235 km s−1.
Paper: Lessons from the High-resolution Spectroscopy of AW UMa and ϵ CrA: Is the Lucy Model Valid?
Authors: Slavek M. Rucinski
Figure caption: A schematic view of ϵ CrA at four orbital phases, as indicated in the figure. The Lucy model with the mass ratio q = 0.13, the fill-out parameter f = 0.25, and the orbital inclination i = 73° was assumed in generating these images. The orbital phases ϕ ≃ 0.57 and ϕ ≃ 0.89 correspond to the orientations when the region of the ESP, shown as a white patch, comes into view or is about to disappear behind the edge of the secondary component. Except for the approximate location in longitude, the geometry and extent of the ESP-causing region in the stellar latitude are unknown and arbitrary in the figure (see the main text).
Paper: Lessons from the High-resolution Spectroscopy of AW UMa and ϵ CrA: Is the Lucy Model Valid?
Authors: Slavek M. Rucinski
Figure caption: A segment of the 2D image formed from individual BFs arranged in orbital phase for the second of the available three observing nights of AW UMa; this is a part of Figure 9 in P1. The radial velocities (the horizontal axis) are in the system centered on the primary component, with the strong rotational profile of the primary component subtracted from the BF profiles for better visibility of the secondary component. The orbital phases (the vertical axis), are expressed as a full count of orbital cycles from the assumed initial epoch, as in P1. The edges of the secondary component profile, as marked by the arrows, E1 and E2, were determined from images similar to this one. Note the strongly nonuniform distribution of intensity across the secondary-component profile, as well as the presence of the bright RV “wisp” that is identified in Figure 7 as ESP #4.
Paper: Lessons from the High-resolution Spectroscopy of AW UMa and ϵ CrA: Is the Lucy Model Valid?
Authors: Slavek M. Rucinski
Figure caption: The orbital-phase evolution of the ESP on the secondary component of ϵ CrA during the transit eclipse. The horizontal axis shows the radial velocity in the heliocentric system while the vertical axis gives the time (increasing upward) expressed in orbital periods (Orb); the eclipse center on the observing night #13 was located at Orb = 3.0 (see P2). The ESP localized at the sub-observer point at ϕ ≃ 0.65 (as described in the text and in Figure 2) may have continued its evolution into the eclipse phases and merged with other spectral-line intensity features. The mass-center velocities of both stars are shown by broken (primary) and dotted (secondary) sine-curve lines.
Paper: Lessons from the High-resolution Spectroscopy of AW UMa and ϵ CrA: Is the Lucy Model Valid?
Authors: Slavek M. Rucinski
Figure caption: Properties of the K. Stȩpień (2009) model follow consideration of the initial stages of the Roche-lobe overflow by the primary. The flow is deflected by the Coriolis force to the side of the secondary, building up a thick belt. The shape of the belt in the orbital plane may be close to an equipotential, but different from that describing the primary component and very likely tighter. As a result, one component slightly overfills and the other slightly underfills the Roche model common equipotential. The figure is for q = 0.10, similar to that of AW UMa and ϵ CrA. The blue line shows the hot gas of the expanding primary, while the secondary component is shown in red. The approximate location of the prominent ESP observed in both binaries is marked in orange. The numbers at the perimeter give the orbital phase positions of an external observer. The arrows at the mass centers represent (in scale) the orbital velocities of the component stars.
Paper: Lessons from the High-resolution Spectroscopy of AW UMa and ϵ CrA: Is the Lucy Model Valid?
Authors: Slavek M. Rucinski
Figure caption: The mass-ratio discrepancy between the photometric, qph, and spectroscopic, qsp, determinations is shown directly versus the spectroscopic value (the upper panel) and normalized by qsp (the lower panel). The data shown in green correspond to qsp determinations from the DDO survey, while those in red are critically updated spectroscopic results listed by P. G. Niarchos & H. W. Duerbeck (1991). See the main text for details.
Paper: Lessons from the High-resolution Spectroscopy of AW UMa and ϵ CrA: Is the Lucy Model Valid?
Authors: Slavek M. Rucinski
Figure caption: The figure shows the AW UMa secondary-component measured velocity extent as vertical lines coded in color by the phase range of the observing night; see the legend in the upper right corner of the figure. The figure is similar to the one for ϵ CrA (Figure 2), where other details are explained. The numbered ESPs observed on individual nights are discussed in the main text.
Paper: Lessons from the High-resolution Spectroscopy of AW UMa and ϵ CrA: Is the Lucy Model Valid?
Authors: Slavek M. Rucinski
Figure caption: The heliocentric orbital velocities of both AW UMa components. The figure corresponds to Figure 4 in P1, with added velocities of the secondary component, derived as the mean values of the secondary profile edges E1 and E2. The scale for the secondary (the right vertical axis) is compressed 10 times relative to that for the primary and is shifted to the same V0. The three observational nights of AW UMa are coded by color: 1–red, 2–green, 3–blue. The continuous lines show the sine curve fits to the data; the dashed line parts correspond to the phase intervals that were not used in the fits because of mutual interference of the wide component profiles. The lower panel shows the half-width of the secondary component profile, W2 = (E2 − E1)/2, which for a standard rotational profile would correspond to the apparent equatorial velocity ﹩{V}_{2}\sin i﹩.
Paper: Lessons from the High-resolution Spectroscopy of AW UMa and ϵ CrA: Is the Lucy Model Valid?
Authors: Slavek M. Rucinski
Figure caption: A schematic view of ϵ CrA at four orbital phases envisaged for the Stȩpień model; compare it with Figure 3 for the Lucy model. The belt surrounding both stars is pictured here with an arbitrary height of  ±0.1 orbital separation units; its temperature is set slightly higher than at the side of the primary for better visibility, while the secondary is cooler. For a simpler visual representation, all shown surfaces follow the equipotentials: the primary is at the inner critical one, and the secondary is at the one with the side radius 0.95 of its inner critical, while the belt is positioned at a point halfway between the inner and outer critical equipotentials.
Paper: Analyses of Multiple Balmer Emission Lines from Accreting Brown Dwarfs and Very Low Mass Stars
Authors: Jun Hashimoto, Yuhiko Aoyama
Figure caption: Out of the 96 data points, 89 were analyzed for line ratios in the Hβ/Hγ vs. Hβ/H8 diagram, with 7 points excluded due to dominant contributions from stellar chromospheric activity. These 89 data points are plotted against the loci of the shock and flow models (Section 3). The data are categorized as follows: 15 points are identified as shock-dominated emission (circles), 55 points as flow-dominated emission (squares), 13 points as fitting both models (crosses), and 6 points as not fitting either model (crosses).
Paper: Analyses of Multiple Balmer Emission Lines from Accreting Brown Dwarfs and Very Low Mass Stars
Authors: Jun Hashimoto, Yuhiko Aoyama
Figure caption: The 96 data points for log(Lacc/L*) are derived from the relationship between accretion luminosity and Hβ line luminosity (J. M. Alcalá et al. 2017), plotted against log(Teff/K). The empirical relationship for chromospheric activity in nonaccreting stars is represented by a dashed line, validated within the range of 3.35 dex ﹩\lesssim {\rm{log}}({T}_{{\rm{eff}}}{\rm{/K}})\lesssim 3.65﹩ dex, following the work of C. F. Manara et al. (2013, 2017). For one object with log(Teff/K) ≈ 3.23 dex, we extrapolate this relationship down to log(Teff/K) = 3.2 dex, indicated by the dotted line. Seven of these 96 data points are identified as dominated by chromospheric emission, as summarized in Table 2 in Appendix A. The 15 solid red circles represent data points categorized as shock-dominated emission in Figure 2, while the two open red circles indicate points identified as chromospheric emission that fall within the shock model loci in Figure 9 in Appendix C.
Paper: Analyses of Multiple Balmer Emission Lines from Accreting Brown Dwarfs and Very Low Mass Stars
Authors: Jun Hashimoto, Yuhiko Aoyama
Figure caption: The 70 data points, categorized as either flow-dominated (blue, 55 points) or shock-dominated (red, 15 points) emissions in Figure 2, are superimposed on diagrams depicting accretion-related physical quantities: (a) freefall velocity (vff) from infinity, (b) accretion luminosity (Lacc), and (c) accretion rate (﹩\dot{M}﹩), all plotted as functions of mass. A black solid line represents the empirical relation ﹩\dot{M}\propto {M}^{2}﹩ across both stellar and substellar regimes (e.g., S. K. Betti et al. 2023). There appears to be no clear boundary between the shock and flow model categories.
Paper: Analyses of Multiple Balmer Emission Lines from Accreting Brown Dwarfs and Very Low Mass Stars
Authors: Jun Hashimoto, Yuhiko Aoyama
Figure caption: The fraction of Hβ line luminosity to accretion luminosity is shown as a function of mass for 70 data points, categorized as either flow-dominated (blue, 55 points) or shock-dominated (yellow in panel (a) and red in panel (b), 15 points) emissions in Figure 2. In panel (a), the accretion luminosity for the 15 shock-dominated data points is derived using stellar scalings (J. M. Alcalá et al. 2017), while in panel (b), it is based on the shock model (Y. Aoyama et al. 2021). For clarity, the error bars are not shown in these panels; however, figures including the error are provided in Figure 11 in Appendix E.
Paper: Analyses of Multiple Balmer Emission Lines from Accreting Brown Dwarfs and Very Low Mass Stars
Authors: Jun Hashimoto, Yuhiko Aoyama
Figure caption: The Hβ spectra are shown for (a) 2MASS J15580252–3736026 with flow-dominated emission, (b) 2MASS J16053215–1933159 with shock-dominated emission, and (c) 2MASS J16083455–2211559 with stellar chromospheric activity dominating. The black dotted lines represent Gaussian profiles with FWHMs of 200.2 km s–1 for (a), 70.7 km s–1 for (b), and 30.6 km s–1 for (c), respectively.
Paper: Analyses of Multiple Balmer Emission Lines from Accreting Brown Dwarfs and Very Low Mass Stars
Authors: Jun Hashimoto, Yuhiko Aoyama
Figure caption: Comparison of mass accretion rates derived using the shock model (﹩{\dot{M}}_{{\rm{A21}}}﹩) and the extrapolation of stellar scalings (﹩{\dot{M}}_{{\rm{A17}}}﹩) for 15 data points located on the shock model loci in Figure 2. The smaller the mass accretion rate, the larger the deviation between the values of ﹩{\dot{M}}_{{\rm{A21}}}﹩ and ﹩{\dot{M}}_{{\rm{A17}}}﹩. Due to the 0.3 dex error in the conversion between LHβ and Lacc,A17 (J. M. Alcalá et al. 2017), the error in the vertical axis is larger.
Paper: Analyses of Multiple Balmer Emission Lines from Accreting Brown Dwarfs and Very Low Mass Stars
Authors: Jun Hashimoto, Yuhiko Aoyama
Figure caption: The line ratio diagrams for (a) Hβ/Hγ vs. Hβ/H6 and (b) Hβ/Hγ vs. Hβ/H8. The loci of line ratios in the shock and flow models are deviated in the left bottom corner as highlighted with dashed lines, by including the H8 hydrogen line at a higher upper level in panel (b). The loci of the shock model appear more discrete than those of the flow model because the parameter grid used in the shock model is coarser. However, this does not impact our overall results.
Paper: Analyses of Multiple Balmer Emission Lines from Accreting Brown Dwarfs and Very Low Mass Stars
Authors: Jun Hashimoto, Yuhiko Aoyama
Figure caption: Comparison of Hγ line profiles for objects observed at multiple epochs, where the emission categories switched between shock and flow models. The spectra were obtained using VLT/X-shooter with a resolution of R ≈ 5400 (55.6 km s–1).
Paper: Analyses of Multiple Balmer Emission Lines from Accreting Brown Dwarfs and Very Low Mass Stars
Authors: Jun Hashimoto, Yuhiko Aoyama
Figure caption: Same as Figure 7, but only including accretion luminosity (upper panels) and mass accretion rate (bottom panels). The accretion luminosity and mass accretion rate of 15 data points, categorized as shock-dominated emissions in Figure 2, are estimated by both the stellar scalings (J. M. Alcalá et al. 2017; yellow) and the shock model (Y. Aoyama et al. 2021; red).
Paper: Analyses of Multiple Balmer Emission Lines from Accreting Brown Dwarfs and Very Low Mass Stars
Authors: Jun Hashimoto, Yuhiko Aoyama
Figure caption: Line ratios of objects with chromospheric-dominated emission selected in Figure 4.
Paper: Analyses of Multiple Balmer Emission Lines from Accreting Brown Dwarfs and Very Low Mass Stars
Authors: Jun Hashimoto, Yuhiko Aoyama
Figure caption: Comparison of the accretion luminosity derived using the shock model (Lacc,A21) with that obtained from the extrapolation of stellar scalings (Lacc,A17) for the 15 data points located on the shock model loci in Figure 2. It is the same as Figure 5, but focuses on accretion luminosity.
Paper: Analyses of Multiple Balmer Emission Lines from Accreting Brown Dwarfs and Very Low Mass Stars
Authors: Jun Hashimoto, Yuhiko Aoyama
Figure caption: Same as Figure 6, but including the error.
Paper: Open-source High-resolution Exoplanet Atmosphere Retrievals with POSEIDON
Authors: Ruizhe Wang, Ryan J. MacDonald, Neale P. Gibson, Nikole K. Lewis
Figure caption: Full corner plots of retrieval results on simulated WASP-121b data. The PT profile of the injection signal is parameterized by six parameters. Blue: Retrieval assuming isothermal atmosphere. Red: Retrieval with six-parameter P-T profile. Upper right: True P-T profile compared to retrieved P-T profile. These results show that high-resolution retrievals are not particularly sensitive to temperature gradients. Runtime: Both about 6 hr on 24 cores.