adsabs/AstroCLIMB huggingface dataset:
Paper: Sculpting of Exoplanetary Systems Driven by a Misaligned Disk and Stellar Oblateness: Origin of Perpendicular Orbits in HD 3167
Authors: Tao Fu, 韬 伏, Yue Wang, 悦 王
Figure caption: Schematic illustration of the dynamical history of the HD 3167 system (not to scale). Consider a primordial near-polar misalignment between the disk (oriented along ﹩{\widehat{{\boldsymbol{l}}}}_{\mathrm{dk}}﹩) and the stellar spin (oriented along ﹩{\widehat{{\boldsymbol{s}}}}_{\star }﹩). Planet b, initially aligned with the disk, transitions to a state in close alignment with the stellar spin during the photoevaporation of the inner region of the disk. The outer three planets (d, c, and e) inherit the orientation of the disk ﹩{\widehat{{\boldsymbol{l}}}}_{\mathrm{dk}}﹩. The outer orbits are strongly coupled, allowing the assumption of their locking together (see the text). The angles indicated in the figure are the observation-constrained values of the HD 3167 system. (See Section 4 for details on the dynamical evolution of HD 3167.)
Paper: Sculpting of Exoplanetary Systems Driven by a Misaligned Disk and Stellar Oblateness: Origin of Perpendicular Orbits in HD 3167
Authors: Tao Fu, 韬 伏, Yue Wang, 悦 王
Figure caption: Dynamical evolution of the HD 3167 system with a near-polar misaligned protoplanetary disk. This simulation uses the full secular model for HD 3167 described in Appendix B. We adopt ψ ⋆,0 = 96° and J 2,0 = 2 × 10−5. Since HD 3167 b with a b ≈ 0.018 au is extremely close to the host star, we further assume that the disk within 0.05 au dissipates with the timescale ﹩{\tau }_{{\rm{v}},\mathrm{in}}^{* }={10}^{3}﹩ yr. All the other parameters are listed in Table 1. From top to bottom, each planet's semi-major axis (along with its pericenter and apocenter distances), obliquity, mutual planetary inclinations, longitude of ascending node, and the dimensionless parameters ϵ obl and ϵ GR are depicted. The observation-constrained values for ψ b⋆, ψ c⋆, and i bc, with their 1σ uncertainties, are also indicated.
Paper: Sculpting of Exoplanetary Systems Driven by a Misaligned Disk and Stellar Oblateness: Origin of Perpendicular Orbits in HD 3167
Authors: Tao Fu, 韬 伏, Yue Wang, 悦 王
Figure caption: Representative examples of orbital evolution for our fiducial system. Three cases of primordial obliquity, ψ ⋆0 = 70° (left), 80° (middle), and 82° (right), are considered. For each case, the panels from top to bottom display the temporal evolution of obliquity, mutual planetary inclination i mut, differences in the ascending node Δh = h b − h c, and the normalized energy integral E and ﹩E\left({L}_{2}\right)﹩. The sky blue dashed curve in the top panel corresponds to the same equilibrium path as depicted in Figure 2(A).
Paper: Sculpting of Exoplanetary Systems Driven by a Misaligned Disk and Stellar Oblateness: Origin of Perpendicular Orbits in HD 3167
Authors: Tao Fu, 韬 伏, Yue Wang, 悦 王
Figure caption: Eccentricity excitation as functions of ϵ obl and ϵ GR. In this figure, we use the full secular model elaborated in Appendix B but without consideration of the disk potential and the decay of stellar oblateness. The initial obliquity is set as ψ ⋆0 = 80°. Four cases of ϵ GR, (a) ϵ GR = 1.0, (b) 5.0, (c) 10.0, and (d) 25.0, are considered. The different ϵ GR values are achieved by changing the semimajor axis of the inner planet a b. We vary the stellar oblateness J 2 in a wide range (as indicated in the top x-axis) to allow ϵ obl to change from 10−2 to 102. The maximum eccentricity e max during the integration up to 5 × 105 yr is recorded for ϵ obl. Except for the parameters specified, all the other simulation parameters are those listed in Table 1.
Paper: Sculpting of Exoplanetary Systems Driven by a Misaligned Disk and Stellar Oblateness: Origin of Perpendicular Orbits in HD 3167
Authors: Tao Fu, 韬 伏, Yue Wang, 悦 王
Figure caption: Domains of different orbit evolution types in the ﹩\left({J}_{\mathrm{2,0}},{\psi }_{\star 0}\right)﹩ parameter space. We vary J 2,0 from 1 × 10−4 to 1 × 10−3 and ψ ⋆0 from 77° to 90°. Four orbit evolution types are distinguished using specific markers, and the final mutual inclination between the inner and outer orbits is color-coded. Except for J 2,0 and ψ ⋆0, all other simulation parameters are those listed in Table 1 for the fiducial system. Note that systems with moderately lower ψ ⋆0 than depicted (e.g., ψ ⋆0 < 78°) tend to experience type 1 evolution, while weaker stellar oblateness (e.g., J 2,0 < 10−4) generally results in type 4 evolution. Results for ψ ⋆0 > 90° are expected to be similar due to the symmetry around ψ ⋆0 = 90°.
Paper: Sculpting of Exoplanetary Systems Driven by a Misaligned Disk and Stellar Oblateness: Origin of Perpendicular Orbits in HD 3167
Authors: Tao Fu, 韬 伏, Yue Wang, 悦 王
Figure caption: Eccentricity stability in the HD 3167 system. This figure is based on simulations using the full secular model for HD 3167 described in Appendix B, excluding effects from disk potential and stellar oblateness decay. The innermost planet’s ϵ GR is calculated to be about 4.67, as indicated by the right vertical dashed line. The green region, where 1 ≲ ϵ obl ≲ ϵ GR, indicates conditions favorable for both large inner–outer mutual inclination and eccentricity stability. The top x-axis shows the values of J 2 corresponding to ϵ obl values on the bottom x-axis.
Paper: JWST Observations of Starbursts: Massive Star Clusters in the Central Starburst of M82
Authors: Rebecca C. Levy, Alberto D. Bolatto, Divakara Mayya, Bolivia Cuevas-Otahola, Elizabeth Tarantino, Martha L. Boyer, Leindert A. Boogaard, Torsten Böker, Serena A. Cronin, Daniel A. Dale, Keaton Donaghue, Kimberly L. Emig, Deanne B. Fisher, Simon C. O. Glover, Rodrigo Herrera-Camus, María J. Jiménez-Donaire, Ralf S. Klessen, Laura Lenkić, Adam K. Leroy, Ilse De Looze, David S. Meier, Elisabeth A. C. Mills, Juergen Ott, Mónica Relaño, Sylvain Veilleux, Vicente Villanueva, Fabian Walter, Paul P. van der Werf
Figure caption: The central 870 pc of M82 seen with NIRCam showing F140M (blue), F250M (green), and F360M (red) on an asinh scale. Insets show zoom-ins to 50 pc × 50 pc regions around some representative massive star cluster candidates. The circles indicate the massive clusters identified in the NIRCam image, where the radius of the circle is the fitted (deconvolved) radius (see Section 3 and Table 1). Blue (pink) circles show clusters which do (do not) overlap with the HST catalog (Y. D. Mayya et al. 2008; see Section 4.1). Clusters 1 and 2 are labeled (see Sections 3.2 and 4.1). The right insets show HST three-color images at the same locations (blue: F435W; green: F555W; red: F814W). The light blue circles in the HST insets show the star clusters identified by Y. D. Mayya et al. (2008) in these selected regions using the same HST data, where the radius of each circle is 5 pc. The pink and blue circles are the same as in the left panels.
Paper: JWST Observations of Starbursts: Massive Star Clusters in the Central Starburst of M82
Authors: Rebecca C. Levy, Alberto D. Bolatto, Divakara Mayya, Bolivia Cuevas-Otahola, Elizabeth Tarantino, Martha L. Boyer, Leindert A. Boogaard, Torsten Böker, Serena A. Cronin, Daniel A. Dale, Keaton Donaghue, Kimberly L. Emig, Deanne B. Fisher, Simon C. O. Glover, Rodrigo Herrera-Camus, María J. Jiménez-Donaire, Ralf S. Klessen, Laura Lenkić, Adam K. Leroy, Ilse De Looze, David S. Meier, Elisabeth A. C. Mills, Juergen Ott, Mónica Relaño, Sylvain Veilleux, Vicente Villanueva, Fabian Walter, Paul P. van der Werf
Figure caption: The cumulative CMFs shown as a fraction of the number of clusters for different assumed ages (colored solid curves). The blue and orange shaded regions on the 0 Myr and 8 Myr curves indicate the approximate range of mass increase due to extinction that scales with M * (dark) or is constant with M * (light); see Section 4.3.1 for details. The gray shaded region shows masses below 104 M ⊙ that are not included in the final catalog. The purple line and shaded region show the power-law fit for the 0 Myr CMF for M * > 104 M ⊙. The dashed orange curve shows the cumulative CMF for the nuclear star clusters identified by Y. D. Mayya et al. (2008) that have an assumed age of 8 Myr.
Paper: JWST Observations of Starbursts: Massive Star Clusters in the Central Starburst of M82
Authors: Rebecca C. Levy, Alberto D. Bolatto, Divakara Mayya, Bolivia Cuevas-Otahola, Elizabeth Tarantino, Martha L. Boyer, Leindert A. Boogaard, Torsten Böker, Serena A. Cronin, Daniel A. Dale, Keaton Donaghue, Kimberly L. Emig, Deanne B. Fisher, Simon C. O. Glover, Rodrigo Herrera-Camus, María J. Jiménez-Donaire, Ralf S. Klessen, Laura Lenkić, Adam K. Leroy, Ilse De Looze, David S. Meier, Elisabeth A. C. Mills, Juergen Ott, Mónica Relaño, Sylvain Veilleux, Vicente Villanueva, Fabian Walter, Paul P. van der Werf
Figure caption: (Bottom left) The mass–radius diagram of the final cluster catalog assuming age = 0 Myr. Typical uncertainties on both quantities are ∼10%. Contours show the 10%, 25%, 50%, and 75% density levels for clusters that overlap with the Y. D. Mayya et al. (2008) catalog and those only detected with NIRCam (pink). The horizontal gray shaded region shows r HWHM less than the F250M HWHM PSF size. The vertical gray shaded region shows M * ≤ 104 M ⊙. The orange arrow shows the shift assuming age = 8 Myr. (Right) A histogram of the deconvolved cluster radii given in Table 1 (gray). The blue and pink histograms show the distribution for clusters that overlap with the Y. D. Mayya et al. (2008) catalog and those that are only detected in NIRCam, respectively. The horizontal axis shows the number of clusters. The gray shaded region is the same as in the mass–radius diagram. (Top) A histogram of the stellar masses for age = 0 Myr given in Table 1 (gray). The pink and blue histograms are as in the right panel. The orange arrow shows the shift assuming age = 8 Myr. The vertical axis shows the number of clusters. The gray shaded region is the same as in the mass–radius diagram.
Paper: JWST Observations of Starbursts: Massive Star Clusters in the Central Starburst of M82
Authors: Rebecca C. Levy, Alberto D. Bolatto, Divakara Mayya, Bolivia Cuevas-Otahola, Elizabeth Tarantino, Martha L. Boyer, Leindert A. Boogaard, Torsten Böker, Serena A. Cronin, Daniel A. Dale, Keaton Donaghue, Kimberly L. Emig, Deanne B. Fisher, Simon C. O. Glover, Rodrigo Herrera-Camus, María J. Jiménez-Donaire, Ralf S. Klessen, Laura Lenkić, Adam K. Leroy, Ilse De Looze, David S. Meier, Elisabeth A. C. Mills, Juergen Ott, Mónica Relaño, Sylvain Veilleux, Vicente Villanueva, Fabian Walter, Paul P. van der Werf
Figure caption: A color–color plot of the clusters (black dots). The contours show the density of the points; contour levels show 10%, 25%, 50%, and 75% density levels for clusters that overlap with the Y. D. Mayya et al. (2008) catalog (blue) and those only detected with NIRCam (pink). The colored tracks show YGGDRASIL SSP models assuming no (f cov = 0; blue) and maximal (f cov = 1; red) contributions from nebular emission. The green arrow shows the slope of the reddening vector with AV  = 3−24 mag marked in increments of 3. The clusters in M82 detected with NIRCam show significant reddening with respect to the dust-free SSP tracks.
Paper: JWST Observations of Starbursts: Massive Star Clusters in the Central Starburst of M82
Authors: Rebecca C. Levy, Alberto D. Bolatto, Divakara Mayya, Bolivia Cuevas-Otahola, Elizabeth Tarantino, Martha L. Boyer, Leindert A. Boogaard, Torsten Böker, Serena A. Cronin, Daniel A. Dale, Keaton Donaghue, Kimberly L. Emig, Deanne B. Fisher, Simon C. O. Glover, Rodrigo Herrera-Camus, María J. Jiménez-Donaire, Ralf S. Klessen, Laura Lenkić, Adam K. Leroy, Ilse De Looze, David S. Meier, Elisabeth A. C. Mills, Juergen Ott, Mónica Relaño, Sylvain Veilleux, Vicente Villanueva, Fabian Walter, Paul P. van der Werf
Figure caption: An example radial profile and Gaussian fit, shown here for Cluster 2 (marked in the upper right JWST inset in Figure 1). The black points show the radial profile extracted from the F250M image, where the error bars reflect the standard deviation in each ring. The blue curve and shaded region show the Gaussian fit to the radial profile and the uncertainties. The vertical blue dashed line indicates the HWHM of the fitted Gaussian. The gray vertical hatched region shows the HWHM of the F250M PSF. The red curve and shaded region show the deconvolved Gaussian fit, and the vertical red dashed line shows the HWHM of the deconvolved Gaussian (reported in Table 1).
Paper: JWST Observations of Starbursts: Massive Star Clusters in the Central Starburst of M82
Authors: Rebecca C. Levy, Alberto D. Bolatto, Divakara Mayya, Bolivia Cuevas-Otahola, Elizabeth Tarantino, Martha L. Boyer, Leindert A. Boogaard, Torsten Böker, Serena A. Cronin, Daniel A. Dale, Keaton Donaghue, Kimberly L. Emig, Deanne B. Fisher, Simon C. O. Glover, Rodrigo Herrera-Camus, María J. Jiménez-Donaire, Ralf S. Klessen, Laura Lenkić, Adam K. Leroy, Ilse De Looze, David S. Meier, Elisabeth A. C. Mills, Juergen Ott, Mónica Relaño, Sylvain Veilleux, Vicente Villanueva, Fabian Walter, Paul P. van der Werf
Figure caption: The progression of ϒ* in the F250M filter with age from the YGGDRASIL models at solar metallicity and maximal nebular emission (see Section 3.3). Key phases of stellar evolution, which govern changes in ϒ*, are marked. See the text in Appendix A for details.
Paper: Ba Enrichment in Gaia MS+WD Binaries: Tracing s-process Element Production
Authors: Param Rekhi, Sagi Ben-Ami, Na'ama Hallakoun, Sahar Shahaf, Silvia Toonen, Hans-Walter Rix
Figure caption: Barium and yttrium abundances as a function of (a) primary mass, (b) eccentricity, (c) period, and (d) Roche-lobe radius of the WD progenitor normalized by the orbital separation. Median error bars are shown in gray. The plotted systems belong to the region M 2 <0.75 and [Fe/H] < −0.08 (see Figure 3 and text for explanation). Although Y-enriched stars extend to high metallicities, we define the s-process domain as per Ba abundances following the existing convention. Roche-lobe radii are computed using the formula given by P. P. Eggleton (1983), with the initial–final mass relation used to obtain progenitor masses taken from J. D. Cummings et al. (2018).
Paper: Ba Enrichment in Gaia MS+WD Binaries: Tracing s-process Element Production
Authors: Param Rekhi, Sagi Ben-Ami, Na'ama Hallakoun, Sahar Shahaf, Silvia Toonen, Hans-Walter Rix
Figure caption: (a) Barium and (b) yttrium abundance as a function of iron abundance of the primary and mass of the secondary (which we assume is a WD). Median error bars are shown in gray. We note that enhanced Ba abundances are found almost exclusively in the region of M 2 < 0.75 and [Fe/H] < −0.08 (bounded in green), with Y-enhanced stars extending to higher metallicities.
Paper: Ba Enrichment in Gaia MS+WD Binaries: Tracing s-process Element Production
Authors: Param Rekhi, Sagi Ben-Ami, Na'ama Hallakoun, Sahar Shahaf, Silvia Toonen, Hans-Walter Rix
Figure caption: Comparison of Ba and Y abundances as a function of iron abundance. Stars are limited to the region of M 2 < 0.75 and [Y/Fe] > 0.25. Y, being a light (first-peak) s-process element, is predicted to be produced more efficiently at higher [Fe/H] values as compared to the heavier (second-peak) Ba (M. Busso et al. 2001). We find [Y/Ba] to increase with [Fe/H], agreeing with the prediction.
Paper: Ba Enrichment in Gaia MS+WD Binaries: Tracing s-process Element Production
Authors: Param Rekhi, Sagi Ben-Ami, Na'ama Hallakoun, Sahar Shahaf, Silvia Toonen, Hans-Walter Rix
Figure caption: Stacked histograms showing the distribution of GALAH derived M 2 in the NCE and RCE samples derived from (a) GALAH DR3 parameters (this work) and (b) Gaia DR3 parameters (S. Shahaf et al. 2024). More than 50% of stars with M 2 ≈ 0.6 M ⊙ formerly classified as RCE get reclassified as NCE systems.
Paper: Ba Enrichment in Gaia MS+WD Binaries: Tracing s-process Element Production
Authors: Param Rekhi, Sagi Ben-Ami, Na'ama Hallakoun, Sahar Shahaf, Silvia Toonen, Hans-Walter Rix
Figure caption: Barium- and yttrium-enriched dwarf stars ([Ba/Fe], [Y/Fe] > 0.25) from the GALAH DR3 survey. The blue and red points have Gaia RUWE values less than and greater than 1.4, respectively. The black dots are the enriched stars in the AMRF-GALAH NCE sample.
Paper: Ba Enrichment in Gaia MS+WD Binaries: Tracing s-process Element Production
Authors: Param Rekhi, Sagi Ben-Ami, Na'ama Hallakoun, Sahar Shahaf, Silvia Toonen, Hans-Walter Rix
Figure caption: Correlations between abundance measurements of s-process elements in the NCE sample. Although GALAH flags high abundances of s-process elements other than Ba and Y as possibly unreliable, we demonstrate here that their abundances are positively correlated with Ba and Y, justifying their use as representative of the s-process in this work.
Paper: Ba Enrichment in Gaia MS+WD Binaries: Tracing s-process Element Production
Authors: Param Rekhi, Sagi Ben-Ami, Na'ama Hallakoun, Sahar Shahaf, Silvia Toonen, Hans-Walter Rix
Figure caption: Distributions of [Ba/Fe] and [Y/Fe] abundances in the AMRF-GALAH sample. The abundance distribution of field stars in GALAH DR3 is shown in dashed blue. A K-S test shows that both the AMRF Ba and Y abundances are inconsistent with the field star distributions with p-values of less than 10−6.
Paper: Cavitating Bubbles in Condensing Gas as a Means of Forming Clumps, Chondrites, and Planetesimals
Authors: E. Chiang
Figure caption: Evolution of bubbles across parameter space. For each of the nonfiducial models, only one of the parameters as listed is changed, while other parameters are left at their fiducial values (namely s = 100 μm, T rad,bkg = 0, P bkg = 0.99P(t = 0), and initial conditions R = 1 km, ﹩\dot{R}=0﹩, T = 2500 K, M = ρ sat(T)V, P = P sat(T), and M solid = 0). Estimated cooling times from Equation (10), shown as vertical dashed lines in the top left panel, predict with varying accuracy the different collapse times (when R ≪ R(t = 0)). Discrepancies between t cool and collapse times are due to evaluating f(τ) in Equation (10) at a single guessed value of τ, when in reality τ varies with time. The longest collapse time, measured in hours, is associated with the R = 100 km bubble, which contains the most mass to cool. The shortest collapse time of 8 s characterizes the bubble initialized with 50% solid particles by mass; these seed solids provide an optical depth near unity at times before the final collapse, thereby maximizing f(τ) (Equation (7)) and by extension the cooling luminosity. The bubble with T rad,bkg > 0 collapses slowest (for this model only, ﹩\dot{R}﹩ in the top right panel is magnified by a factor of 10 for ease of viewing) and has its short-period pressure oscillations damped away at late times.
Paper: Cavitating Bubbles in Condensing Gas as a Means of Forming Clumps, Chondrites, and Planetesimals
Authors: E. Chiang
Figure caption: Adding hydrogen, a noncondensible gas, to the bubble maintains a minimum pressure and prevents wholesale collapse. Aside from the H2 fraction, all other parameters of the models shown here are the same as in the fiducial model. Even small hydrogen mass fractions can effectively stabilize the bubble. The initial hydrogen partial pressure ﹩{P}_{{{\rm{H}}}_{2}}﹩ is larger than the initial saturation vapor pressure of silicates by the ratio of mean molecular weights ﹩\mu /{\mu }_{{{\rm{H}}}_{2}}\simeq 15﹩. When ﹩{M}_{{{\rm{H}}}_{2}}/M=0.01﹩, the hydrogen partial pressure ﹩{P}_{{{\rm{H}}}_{2}}﹩ sits just a factor of 10 below the background pressure P bkg (assumed to be nearly equal to the bubble’s initial total pressure). This factor of 10 is recovered, and pressure equilibrium with the background achieved, once the bubble shrinks by a factor of ∼2 from the loss of silicate vapor pressure (top panels). At this point all of the silicate is condensed; the dust particles, now residing in inert H2, are free to radiatively cool to the background radiation temperature (here zero), further shrinking the bubble but more slowly compared to earlier. From dT/dt ∝ R 2 T 4/τ (for optically thick bubbles; see bottom right panel), τ ∝ R, and T/V = constant (constant pressure maintained by the background), it follows that T ∝ t −3/10 and R ∝ t −1/10. We have verified that secular cooling and shrinkage halt once T reaches T rad,bkg > 0 (data not shown).
Paper: Cavitating Bubbles in Condensing Gas as a Means of Forming Clumps, Chondrites, and Planetesimals
Authors: E. Chiang
Figure caption: Evolution of our fiducial bubble. The middle panel is the same as the top panel, zoomed in to show small variations in temperature, pressure, and density (gas density ρ gas is distinct from the total density ρ in solids+gas). The variations reflect fast acoustic oscillations and slower secular cooling by radiation. Initial conditions are listed in the top panel; additional inputs include the initial bubble velocity ﹩\dot{R}(t=0)﹩ and the background radiation temperature T rad,bkg, both zero here. The bubble is seeded with a percent-level overpressure relative to the background (dotted orange line in middle panel) that causes it to first expand (solid black line), cool (dashed blue line), and condense (bottom panel showing dust optical depth τ and fractional mass in solids M solid/M). Radiation from this dust cools the bubble, and subsequent temperature oscillations trend ever so slightly downward. The corresponding secular decline in pressure causes the bubble to shrink. The smaller volume leads to still more condensation undersaturated conditions (Equation (3)). When the bubble has nearly fully condensed (t ≳ 14 s), the order-unity pressure imbalance between bubble and background results in a final implosion. The cooling time given by Equation (10) is marked by a vertical dashed line.
Paper: Cavitating Bubbles in Condensing Gas as a Means of Forming Clumps, Chondrites, and Planetesimals
Authors: E. Chiang
Figure caption: Trajectories of solid particles of varying sizes s  c  (distinct from s, the size of the particles controlling opacity) within six of the model bubbles shown in Figure 2 (omitting the cold, initial T = 2000 K model). The particles’ radial locations R  c  are normalized to the bubble radius R. Particles are initially placed at either R  c /R = 0.25 or 0.75; to more easily distinguish between curves, some have been offset in the ordinate, or scaled in time by order-unity factors. The largest particles in each panel fail to be entrained by bubble gas and are stranded at R  c /R > 1 while the bubble collapses. Smaller particles stay inside and in some cases are propelled to the center somewhat faster than gas.
Paper: The MIRI Exoplanets Orbiting White dwarfs (MEOW) Survey: Mid-infrared Excess Reveals a Giant Planet Candidate around a Nearby White Dwarf
Authors: Mary Anne Limbach, Andrew Vanderburg, Alexander Venner, Simon Blouin, Kevin B. Stevenson, Ryan J. MacDonald, Sydney Jenkins, Rachel Bowens-Rubin, Melinda Soares-Furtado, Caroline Morley, Markus Janson, John Debes, Siyi Xu, Evangelia Kleisioti, Matthew Kenworthy, Paul Butler, Jeffrey D. Crane, Dave Osip, Stephen Shectman, Johanna Teske
Figure caption: The measured IR excess or deficit divided by the precision of our measurement (y-axis) in all three spectral bands (x-axis) for a subset of MEOW white dwarfs. The orange line represents WD 0310–688, which shows a highly significant IR excess. The purple line corresponds to another white dwarf with an excess possibly due to an exoplanet but detected with less significance than WD 0310–688. The green lines indicate no significant excess, aligning generally with our models, although they are systematically low by 1%–2% at 7.7 μm and 21 μm. The white dwarf that is notably low at 21 μm is likely affected by a nearby bright red galaxy contaminating the background flux measurement.