# The Radial Distribution and Excitation of H2 around Young Stars in the
$HST$-ULLYSES Survey

Kevin France Laboratory for Atmospheric and Space Physics, University of
Colorado Boulder, Boulder, CO 80303  Nicole Arulanantham Space Telescope
Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA  Erin
Maloney  Laboratory for Atmospheric and Space Physics, University of Colorado
Boulder, Boulder, CO 80303  P. Wilson Cauley  Laboratory for Atmospheric and
Space Physics, University of Colorado Boulder, Boulder, CO 80303  P. Ábrahám
Konkoly Observatory, Research Centre for Astronomy and Earth Sciences, Eötvös
Loránd Research Network (ELKH), Konkoly-Thege Miklós út 15-17, 1121 Budapest,
Hungary  CSFK, MTA Centre of Excellence, Budapest, Konkoly Thege Miklós út
15-17., H-1121, Hungary  ELTE Eötvös Loránd University, Institute of Physics,
Pázmány Péter sétány 1/A, H-1117 Budapest, Hungary  Juan M. Alcalá  INAF-
Osservatorio Astronomico di Capodimonte, via Moiariello 16, 80131 Napoli,
Italy  Justyn Campbell-White European Southern Observatory, Karl-
Schwarzschild-Strasse 2, 85748 Garching bei München, Germany  Eleonora
Fiorellino INAF-Osservatorio Astronomico di Capodimonte, via Moiariello 16,
80131 Napoli, Italy  Gregory J. Herczeg Kavli Institute for Astronomy and
Astrophysics, Peking University, Beijing 100871, China  Department of
Astronomy, Peking University, Beijing 100871, China  Brunella Nisini INAF -
Osservtorio Astronomico di Roma. Via di Frascati 33, 00078 Monte Porzio
Catone, Italy  Miguel Vioque Joint ALMA Observatory, Alonso de Córdova 3107,
Vitacura, Santiago 763-0355, Chile  National Radio Astronomy Observatory, 520
Edgemont Road, Charlottesville, VA 22903, USA Kevin France
<EMAIL_ADDRESS>
###### Abstract

The spatial distribution and evolution of gas in the inner 10 au of
protoplanetary disks form the basis for estimating the initial conditions of
planet formation. Among the most important constraints derived from
spectroscopic observations of the inner disk are the radial distributions of
the major gas phase constituents, how the properties of the gas change with
inner disk dust evolution, and how chemical abundances and excitation
conditions are influenced by the high-energy radiation from the central star.
We present a survey of the radial distribution, excitation, and evolution of
inner disk molecular hydrogen (H2) obtained as part of the $HST$/ULLYSES
program. We analyze far-ultraviolet spectroscopy of 71 (63 accreting) pre-main
sequence systems in the ULLYSES DR5 release to characterize the H2 emission
lines, H2 dissociation continuum emission, and major photochemical/disk
evolution driving UV emissions (Ly$\alpha$, UV continuum, and C IV). We use
the widths of the H2 emission lines to show that most fluorescent H2 arises
between 0.1 – 1.4 au from the parent star, and show positive correlations of
the average emitting radius with the accretion luminosity and with the dust
disk mass. We find a strong correlation between H2 dissociation emission and
both the accretion-dominated Ly$\alpha$ luminosity and the inner disk dust
clearing, painting a picture where water molecules in the inner 3 au are
exposed to and dissociated by strong Ly$\alpha$ emission as the opacity of the
inner disk declines with time.

††journal: AAS journals

##  1 INTRODUCTION

The circumstellar environments in which planets form are dynamic and volatile
places. Mass accretion onto the central stars is active and rapid stellar
rotation drives strong magnetic fields. These processes power high levels of
UV and X-ray emission, as well as temporal variability that alters the
physical state and longevity of the protoplanetary disk. The stellar high-
energy radiation (UV and X-ray) controls circumstellar disk dispersal
(Alexander et al. 2014; Pascucci et al. 2022), strongly influencing the final
stages of planet formation. Young stars complete their mass assembly at the
same time as protoplanets are forming around these objects (1 – 10 Myr; Haisch
et al. 2001; Fedele et al. 2010; Ribas et al. 2014), therefore, there is
strong competition between dynamical disk clearing through forming
protoplanets (e.g., Rice et al. 2003; Dodson-Robinson & Salyk 2011; van der
Marel et al. 2018) and disk dispersal through photoevaporation and disk winds,
which depend on the strength and shape of the pre-main sequence far-
ultraviolet (FUV; 912 – 1800 Å) radiation field (Alexander et al. 2014;
Ercolano & Owen 2016; Ercolano & Pascucci 2017; Wang & Goodman 2017). The
high-energy radiation produced in the accretion region and surrounding shock-
heated gas is responsible for the high-levels of UV continuum and line
radiation from Classical T Tauri Stars (CTTSs; Ingleby et al. 2011; France et
al. 2014b; Schneider et al. 2020).

Figure 1: Example $HST$-COS spectra of four representative CTTSs from the
ULLYSES survey. The stellar spectra are color-coded and labeled in the legend
in the lower left; offsets were applied to Sz 69 (5 $\times$ 10-15 erg cm-2
s-1 Å-1), RXJ1556.1-3655 (the flux was divided by 2 and an offset of 5
$\times$ 10-15 erg cm-2 s-1 Å-1 was applied), and XX Cha (20 $\times$ 10-15
erg cm-2 s-1 Å-1). A 3 resolution element (21 pixel) boxcar smooth has been
applied for plotting the full spectral range. The wavelengths of prominent
emission lines of H2 (green diamonds, separated by lines that are primarily
recorded in the G130M or G160M modes) and atomic lines (red circles) are
indicated. The peak wavelength region of the 1600 Å Bump emission from
Ly$\alpha$ dissociated H2 (originating from a Ly$\alpha$ dissociated H2O
parent population) is indicated by the gray shaded region. The $HST$-COS
observations of CVSO 176 used a single central wavelength setting which
results in a $\approx$ 12 – 15 Å gap in the spectra where the COS detector
segments a and b meet.

Current studies of the accretion of circumstellar disk material onto young
stars explore the competing processes of magnetorotational instability (MRI)
and magnetohydrodynamic (MHD) disk winds for removing excess angular momentum
of from the disk and allowing material to accrete onto the central protostar.
In the MRI scenario, gas within $\sim$ 0.1 au is ionized by the central star
and entrained in the stellar magnetic field (Balbus & Hawley 1991), while the
MHD disk wind picture proposes that outflowing gas from the disk surface at
radii out to tens of au (see, e.g., Lesur 2021 for a recent review) removes
angular momentum from the disk. Ultimately, material is funneled onto the
central star along field lines where it creates an accretion shock at
(possibly) several contact points with the pre-main sequence photosphere
(Calvet & Gullbring 1998; Muzerolle et al. 2001; Hartmann et al. 2016; Pittman
et al. 2022). The final phase of mass accretion (stellar ages 2 – 10 Myr) is
particularly important as accretion luminosity and stellar activity catalyze
the dispersal and chemistry of the protoplanetary disk, impacting the
timescales for planet formation and the accretion of protoplanetary
atmospheres (Wang & Goodman 2017).

The high-energy radiation from the accretion region is also responsible for
ionizing and photodissociating inner disk atoms and molecules, catalyzing disk
surface chemistry, and driving many of the observable signatures of the inner
regions of the disk. Molecular gas emission and absorption originating inside
of 10 au provide our best means of estimating the conditions at the radii
where gas giant and rocky planet cores are forming and accreting their nascent
atmospheres. Surveys of molecular emission from the inner few au provide
constraints on the radial distribution, temperature, and composition of
planet-forming inner disks. Surveys of mid-IR emission from CO (Salyk et al.
2009; Brown et al. 2013; Banzatti & Pontoppidan 2015; Banzatti et al. 2022),
H2O and organic molecules (Pontoppidan et al. 2010; Carr & Najita 2011; Salyk
et al. 2011b; Banzatti et al. 2023), UV emission and absorption of H2 and CO
(Herczeg et al. 2004; France et al. 2011; Schindhelm et al. 2012a; France et
al. 2014a; Arulanantham et al. 2021), and spectrally/spatially-resolved near-
IR observations (Pontoppidan et al. 2011; Carmona et al. 2011; Brittain et al.
2015) have placed constraints on the relative abundance ratios of H2, CO, and
H2O, the evolution of the inner gas disk radius, and the excitation conditions
of the molecular gas. The inventory of disk-molecules and photochemical
studies of the planet-forming environments around young stars is expected to
expand rapidly in the era of $JWST$ (Öberg & Bergin 2021; Grant et al. 2022;
Kóspál et al. 2023).

H2 is the primary mass component of gas-rich disks, however, the lack of
intrinsic dipole moment and large energy spacing between the rotational levels
have traditionally made direct detection of the rovibrational emission lines
of the molecule challenging to study in a large number of disks (although see,
e.g., Carmona et al. 2011; Gangi et al. 2020; Kóspál et al. 2023). Conversely,
the electronic excitation spectrum of H2 is among the brightest features in UV
spectra of pre-main-sequence stars (France et al. 2012a), and can serve as a
diagnostic of weak mass accretion when traditional accretion diagnostics
(e.g., H$\alpha$, NUV Balmer continuum) do not provide clear evidence of
active accretion (Ingleby et al. 2011; Alcalá et al. 2019).

The molecular tracers observed at UV wavelengths are primarily excited by UV
fluorescence and dissociation. Fluorescent H2 (Herczeg et al. 2002) emission
is detected in the circumstellar environments of all accreting protostars
surveyed to date (e.g., France et al. 2012a) and is a sensitive probe of
remnant gas in disks with evolved dust opacity (Hoadley et al. 2015).
Fluorescent CO emission is observed in a fraction of disks, with an observed
bias towards disks that show dust clearing in their mid-IR spectral slopes
(Schindhelm et al. 2012a; Arulanantham et al. 2021). H2 dissociation continuum
emission (the “1600 Å Bump”), where Ly$\alpha$ photons dissociate water and
the resultant, highly non-thermal H2 dissociation products, is observed in
$\sim$all dust-evolved disks (France et al. 2017), however, only approximately
a third of “full” or “primordial” disks display this feature. Together, the
fluorescent H2 line emission and continuum bump trace both the radial
distributions and UV irradiation of hot molecular gas in the inner disks.

In this paper, we present a new survey of the H2 features observed in the UV
spectra of CTTSs, taking advantage of the Ultraviolet Legacy Library of Young
Stars as Essential Standards (ULLYSES) program. ULLYSES obtained almost 500
orbits of ultraviolet spectroscopic observations with the Hubble Space
Telescope ($HST$), targeting approximately 71 new accreting and non-accreting
young stars. We analyze the ULLYSES Data Release 5 (DR5) to provide the
largest spectral catalog of H2 line and continuum features obtained to date to
better understand the radial distribution and evolution of the inner disk
molecules around young stars.

This paper is laid out as follows: Section 2 presents an overview of the
ULLYSES FUV spectra analyzed here and Section 3 presents the emission line
measurements and a publicly-available data catalog of the FUV line parameters.
Section 3 concludes with the derivation of the Ly$\alpha$ luminosity, average
H2 fluorescent emitting radius, and the properties of the FUV continuum
radiation. In Section 4, we discuss how these results suggest a physical
picture where stellar/accretion photons have greater penetration depth and
drive greater molecular dissociation as the dust opacity of the inner disk
declines. We conclude with a brief summary in Section 5.

Figure 2: Individual H2 $B$${}^{1}\Sigma^{+}_{u}$ – $X$${}^{1}\Sigma^{+}_{g}$
(1 – 7) R(3) ($\lambda_{lab}$ = 1489.57 Å) emission lines for the four example
targets from Figure 1. The spectra are shifted to their rest wavelength for
comparison and normalized by the average flux in the peak ($\pm$ 10 km s-1) of
the emission line and offset by 1.0. The $HST$-COS line-spread-function (LSF)
is shown as the dotted green line; all emission lines are well-resolved in the
COS data, indicating kinematic line-broadening.

##  2 Archival Far-Ultraviolet Spectra from ULLYSES

We analyze FUV spectra obtained as part of the 487 $HST$ orbits in the low-
mass star portion of ULLYSES (see Roman-Duval et al. 2020 for a program
overview). The observations include data acquired through Data Release
5111https://ullyses.stsci.edu/ullyses-download.html with the addition of five
additional targets observed in Summer 2022. The final sample (Table 1) is
primarily composed of accreting T Tauri stars observed as part of the ULLYSES
program, however, the DR5 catalog includes a number of stars from previous
young star observing programs with similar observational plans (see, e.g.,
Ardila et al. 2013; France et al. 2017). The final source list analyzed here
includes 63 accreting and 8 non-accreting sources.

The target sample presented in ULLYSES DR5 were obtained between late-2009 and
mid-2022 with the $HST$-Cosmic Origins Spectrograph (Green et al. 2012; COS).
As the data collection strategy and analysis of FUV spectra of accreting
protostars has been discussed extensively in previous works (Roman-Duval et
al. 2020; France et al. 2012a), we present a high-level overview of the
dataset here with an emphasis on the FUV spectral line catalog creation in
Section 3. The FUV spectral reduction and analysis follow the general
description from the ODYSSEUS program overview paper (Espaillat et al. 2022),
and we refer the reader to that work for an example of the FUV spectra
analyzed in the context of the larger ODYSSEUS panchromatic analysis program.
We also employ mid-IR spectral slopes (12 – 22$\mu$m) from $WISE$ photometry
(Cutri et al. 2013). The $WISE$ data was downloaded from the NASA IPAC
Infrared Science Archive and $WISE$ observations with data quality flags
(cc_flags $\neq$ 0) were excluded from the analysis. In addition, visual
inspection of the $WISE$ images indicated potential field crowding or extended
emission in CVSO-90, CVSO-107, CVSO-176, RECX-15 (ET Cha, echa j0843.3-7915),
and CS Cha. Of these, we note that CS Cha hosts a well-studied transition disk
(Espaillat et al. 2007) providing confidence in the large $WISE$ W3 – W4 color
for this source.

The $HST$ observations were made using the COS G130M and G160M grating modes.
Most observing programs employed multiple central wavelength settings at
several focal-plane split positions to create continuous FUV spectra from
$\sim$ 1140 – 1760 Å and mitigate the effects of fixed pattern noise222We note
that some of the ULLYSES observations used a single central wavelength setting
which results in $\approx$ 12 – 15 Å gaps in the spectra where the COS
detector segments a and b meet.. These modes provide a point-source resolution
between $\Delta$ $v$ $\approx$ 17 – 22 km s-1 with 6 – 7 pixels per resolution
element (Osterman et al. 2011). The total FUV exposure times were between two
and sixteen orbits per target, depending on the intrinsic luminosity and the
interstellar plus circumstellar reddening on the sightline.

We present an example of four representative T Tauri star spectra from ULLYSES
in Figure 1 (XX Cha, RXJ1556.1-3655, Sz 69, and CVSO 176). These stars come
from a variety of star-forming regions sampled by ULLYSES (Chamaeleon I,
Corona Australis, Lupus, and Orion OB1) and the locations of prominent
molecular and atomic features in the spectra are shown. As this work focuses
primarily on the line and continuum emission from H2, Figures 2 and 3 present
representative samples of these features. Figure 2 shows a representative,
isolated, fluorescent emission line of H2 ($B$${}^{1}\Sigma^{+}_{u}$ –
$X$${}^{1}\Sigma^{+}_{g}$ (1 – 7) R(3); $\lambda_{lab}$ = 1489.57 Å) for the
four stars shown in Figure 1. Figure 3 presents the range of H2 continuum
spectra observed in ULLYSES DR5: the top panel shows the underlying FUV
spectral continuum (CVSO-109) showing no molecular emission (see also
Espaillat et al. 2022) while the middle panel shows the same spectral
extraction from a source with a strong molecular continuum feature (“1600 Å
Bump”, France et al. 2017; DM Tau). The bottom panel of Figure 3 shows the
extracted H2 molecular continuum from DM Tau (see Section 3.3).

Figure 3: Example FUV continuum spectra for two targets in the ULLYSES DR5,
CVSO 109 (top panel; Espaillat et al. 2022) and DM Tau (center panel). The
underlying accretion continuum (fit with a polynomial; blue dashed line) is
removed and the residual molecular continuum is shown as the black circle data
points in the bottom panel. The molecular continuum fit is parameterized by a
Gaussian + polynomial (France et al. 2017) and is shown in red. Molecular
continuum fit parameters are shown in the legend; the log10(luminosity) of the
FUV continuum and the 1600 Å Bump, in units of erg s-1, are presented as
L(cont) and L(bump), respectively

.

##  3 FUV Spectral Line Data Reduction and Analysis

We analyzed the brightest FUV lines in the COS spectra: five progressions of
fluorescent H2 ([1,7], [1,4], [0,1], [0,2], and [2,12]) and hot gas emission
from Si III (1206 Å), N V (1238 Å and 1243 Å), Si IV (1394Å and 1403Å), and C
IV (1548Å and 1550Å). The H2 lines are marked in Figure 1, depending on
whether they were recorded on the G130M (open diamonds) or G160M (filled
diamonds) grating and the hot gas lines are marked with red circles. The
spectra and tables of line flux measurements are publicly available in
machine-readable format333https://lasp.colorado.edu/home/cusp/outflows-and-
disks/. For each star, we take the high-level processed data from the ULLYSES
*cspec.fits files and extract individual emission line flux and error regions,
which are saved as line-specific arrays. For hot gas doublets (e.g., C IV, N
V, and Si IV), the lines are extracted separately into the blue and red
doublet components. Individual extracted emission line region and tables of
line measurements for the ensemble of stars are available on the web portal.

Each array is processed with a 5-pixel boxcar smoothing algorithm to suppress
noise without compromising significant kinematic information (all of the
emission lines studied here are well-resolved see, e.g., France et al. 2012b;
Figure 2). The $HST$-COS line spread function contributes $\sim$20% for lines
narrower than 40 km s-1 and less for broader lines. For each emission line
region, background continuum flux is calculated from an immediately adjacent
spectral region then subtracted out from the entire flux array. In Table 2, we
present the integrated fluxes measured over the interval $\pm$ 200 km s-1.

The full width at half maximum (FWHM) was calculated non-parametrically to
account for the various non-Gaussian shapes of some of the emission lines
(Ardila et al. 2013). The brightest 7 pixels (approximately one resolution
element) in an emission line region are averaged to find the maximum flux,
Fmax. The line center is defined as the average wavelength within the 7 pixel
Fmax region. The half maximum value, $F_{max/2}$ is calculated as

$F_{max/2}=(F_{max}-F_{cont})/2+F_{cont}$ (1)

where $F_{cont}$ is the value of the continuum emission near the line of
interest. The $\pm$ half-width points are the wavelengths corresponding to
Fmax/2; the difference of these half-width points are the FWHMs presented in
Table 2. For the hot gas doublets the FWHM was only calculated on the blue
component of the doublet because these doublets are governed by the same
physical creation mechanism and the blue components of the doublet are
stronger (brighter). Uncertainties on the FWHM are computed using the same
method, with the $\pm$1 $\sigma$ flux error added to the flux values to
evaluate the minimum and maximum values achievable within the photometric flux
uncertainty.

In certain instances (approximately 5% of line profile measurements), the FWHM
calculation returned values that did not match a qualitative inspection of the
data, and for these lines, we measured the line profiles following the method
described by France et al. (2012b), where a Gaussian line shape is forward-
modeled and compared to the data using a Levenberg-Marquardt least-squares
minimization to return the best-fit emission line parameters. This “patch” was
tested against approximately a dozen emission lines with good non-parametric
results, and the agreement in FWHM and flux was found to be within 10-20% in
all cases. The emission line measurement files are stored as comma separated
value arrays and are publicly available 444
https://lasp.colorado.edu/home/cusp/outflows-and-disks/  . Each data file
includes the emission line IDs, the peak wavelength of each line, the
integrated flux of the line, the error on that integrated flux, the FWHM of
the line, the error on that FWHM and a data quality flag. The data quality
flags note special circumstances or indicate when there is no emission line.
In the following subsections, we describe how these basic data products were
used to derive system parameters of the ULLYSES targets (§3.1 and §3.2) and
how the FUV continua were extracted from the $HST$-COS datasets (§3.3).

###  3.1 H2 Progression and Ly$\alpha$ Luminosity from the ULLYSES Sample

The fluorescent H2 lines observed in the CTTS sample can be used to constrain
the spatial distribution of H2 in the circumstellar environment. For our line-
profile analysis, we focus on the measurement of four progressions
([$v^{{}^{\prime}}$,$J^{{}^{\prime}}$] = [1,7], [1,4], [0,1], and [0,2]).
These progressions are pumped through the (1 – 2)R(6) $\lambda_{lab}$1215.73
Å, (1 – 2)P(5) $\lambda_{lab}$1216.07 Å, (0 – 2)R(0) $\lambda_{lab}$1217.21 Å,
and (0 – 2)R(1) $\lambda_{lab}$1217.64 Å absorbing transitions,
respectively555The quantum numbers $v$ and $J$ denote the vibrational and
rotational quantum numbers in the ground electronic state
($X$${}^{1}\Sigma^{+}_{g}$), the numbers $v^{{}^{\prime}}$ and
$J^{{}^{\prime}}$ characterize the H2 in the excited electronic state
($B$${}^{1}\Sigma^{+}_{u}$), and the numbers $v^{{}^{\prime\prime}}$ and
$J^{{}^{\prime\prime}}$ are the rovibrational levels in the electronic ground
state following the fluorescent emission. Absorption lines are described by
($v^{{}^{\prime}}$ – $v$) and emission lines by ($v^{{}^{\prime}}$ –
$v^{{}^{\prime\prime}}$).. The absorbing transitions are within +14 – +487 km
s-1 of Ly$\alpha$ line-center. We target the brightest two, non-blended
emission lines from each progression for analysis here. The H2 emission line
centroids are broadly consistent with the stellar radial velocities within the
15 km s-1 accuracy of the $HST$-COS wavelength solution, and we will explore
multiple emission line components (e.g., Arulanantham et al. 2018) and their
kinematic signatures in a future work.

The emission line fluxes are corrected for interstellar reddening using the AV
values given in Table 1. It has been shown that the dust properties of the
Taurus star-forming region differ from the Milky Way average extinction curve
(Calvet et al. 2004). However, the shape of the extinction curve for star-
forming regions and individual targets remains highly uncertain (McJunkin et
al. 2016), and the AV values derived for individual target carry significant
uncertainties (Pittman et al. 2022; Carvalho & Hillenbrand 2022). Both of
these issues are exacerbated by the strong wavelength-dependence of the dust
attenuation curve, where the reddening correction at 1500 Å has a $\sim$ 6x
larger amplitude relative to 6500 Å for AV = 1. In light of these
uncertainties, we adopt a typical interstellar curve Cardelli et al. (1989)
and an $R_{V}$ = 3.1. If, on the other hand, we would have adopted a standard
interstellar curve with $R_{V}$ = 5.5 (for AV = 1), the resultant fluxes would
be between 20 – 40% of the values reported here.

The total flux from a progression $m$ is given by

$F_{m}(\textup{H${}_{2}$})=\frac{1}{2}\sum\left(\frac{F_{mn}}{B_{mn}}\right)$
(2)

where $F_{mn}$ is the reddening corrected, integrated H2 emission line flux
from rovibrational state $m$ (= [$v^{{}^{\prime}}$,$J^{{}^{\prime}}$]) in the
$B$${}^{1}\Sigma^{+}_{u}$ electronic state to $n$ (=
[$v^{{}^{\prime\prime}}$,$J^{{}^{\prime\prime}}$]) in the ground electronic
state, $X$${}^{1}\Sigma^{+}_{g}$. $B_{mn}$ is the branching ratio between
these two states, and 2 is the number of emission lines measured from a given
progression. The measurement errors are typically small, so we take the flux
error to be the average error of the individual measurements of $F_{m}$(H2).
The dominant systematic error on the measured H2 flux is the correction for
interstellar reddening; we do not attempt to account for this uncertainty in
the flux and luminosity errors presented below. The total progression
luminosity is then $L_{m}$(H2) = (4$\pi$$d^{2}$)$F_{m}$(H2). In Table 2, we
present reddening-corrected luminosities for the [1,4] and [0,1] progressions.
The total H2 luminosity is the sum of the individual progression luminosities.
Because we select a finite number of progressions to analyze, this treatment
underestimates the total H2 fluorescent luminosity by an amount that depends
on the temperature and spatial distribution of the absorbing H2, as well as
the shape of the Ly$\alpha$ pumping spectrum. As these four progressions are
the brightest in essentially all CTTSs, we consider this to be a 20 – 30%
underestimate of the total fluorescent H2 emission (France et al. 2012a).

Figure 4: Projected Ly$\alpha$ luminosities for all ULLYSES DR5 stars (black
diamonds) as a function of total fluorescent H2 luminosity. We created a
power-law fit to the 14 stars with high-quality, multi-progression H2 emission
line data and Ly$\alpha$ reconstructions from France et al. (2012a), shown as
orange circles on the plot. This fit was used to predict the Ly$\alpha$
luminosity of the ULLYSES stars based on their observed F(H2).  Figure 5:
Comparison of the C IV and Ly$\alpha$ luminosities of the ULLYSES stars (black
diamonds) and previous $HST$-COS survey of CTTS (France et al. 2012b; orange
circles). The plot demonstrates the correlation between Ly$\alpha$ (as derived
in Section 3.1) and accretion-dominated emission from C IV. This is consistent
with the broad, bright Ly$\alpha$ profiles of CTTS (Schindhelm et al. 2012b;
Arulanantham et al. 2023) being powered by mass accretion (e.g., France et al.
2014b).

The fluorescent emission lines of H2 can be used to reconstruct the bright
Ly$\alpha$ emission line that is the primary excitation source of this gas
(Wood et al. 2002; Herczeg et al. 2004; Schindhelm et al. 2012b). In order for
this process to produce a high-fidelity reconstruction, a sufficient number of
H2 progressions, covering the majority of the Ly$\alpha$ pumping profile, must
be detected at moderate to high S/N (e.g., Schindhelm et al. 2012b;
Arulanantham et al. 2018). Ly$\alpha$ reconstructions with a small number of
low-S/N progressions leave the shape and integrated flux of the Ly$\alpha$
line poorly constrained (Espaillat et al. 2022). Of the 63 accreting sources
analyzed here, the majority are only detected in 3 or 4 H2 progressions – 54
detections in [1,4], 41 detections in each [1,7] and [0,1], 33 detections in
[0,2], and only 3 detections in the [2,12] progression that is often used to
constrain the wings of the Ly$\alpha$ line. Therefore, a full Ly$\alpha$
reconstruction cannot be practically applied to the full ULLYSES sample.

In order to provide a uniform estimate for the Ly$\alpha$ luminosity for
ULLYSES, we fit F(H2) vs. F(Ly$\alpha$) relation using the 14 higher S/N
targets reconstructed by Schindhelm et al. (2012b) and H2 fluorescence data
from France et al. (2012a). These data are shown as the orange points in
Figure 4. We extrapolated the ULLYSES measurements of L(H2) described above
onto this relation to estimate the Ly$\alpha$ fluxes for all 63 accreting
stars in our sample:

$log_{10}(F(Ly\alpha))=m_{Ly\alpha}(log_{10}(F(H_{2})))+b_{Ly\alpha}$ (3)

where the best-fit coefficients are $m_{Ly\alpha}$ = 0.791 $\pm$ 0.033 and
$b_{Ly\alpha}$ = $-$1.297 $\pm$ 0.400. Propagating the uncertainties on the
fit coefficients of the F(H2) vs. F(Ly$\alpha$) relation indicates a maximum
uncertainty on the associated Ly$\alpha$ luminosities of $\lesssim$40%. The
small dispersion seen in the extrapolated Ly$\alpha$ luminosities is due to
the conversion of the flux-flux relationship from France et al. (2012a) to
luminosity. We caution that while the original Ly$\alpha$-to-H2 relationship
from France et al. (2012a) spans approximately 1.5 orders of magnitude in
L(H2), we are extrapolating that relationship to over 3 orders of magnitude
for the present sample. Figure 5 illustrates the comparison of the estimated
Ly$\alpha$ to C IV luminosity between the present work and the (France et al.
2012a) sample, showing reasonable agreement in the range $L$(C IV) $>$ 1029
erg s-1 with larger spread in the lower C IV luminosity region of the ULLYSES
sample. The spread in Ly$\alpha$ values at $L$(C IV) $<$ 1029 erg s-1 is
driven by scatter in the relationship between C IV and H2. Similarly, we a
moderate positive correlation between L(Ly$\alpha$) and the photometric
luminosity (L∗; from Manara et al. 2021, 2022): the Spearman rank correlation
coefficient, $\rho$, and the likelihood that this correlation is consistent
with a random distribution, n, for the L(Ly$\alpha$) vs. L∗ relation are
$\rho$ = 0.498 and n = 2.3 $\times$ 10-4.

###  3.2 Radial Distribution of H2 Fluorescent Emission

Kinematic broadening dominates the observed H2 line profiles. The thermal
broadening of the emission lines is approximately 4.5 km s-1 at the nominal
2500 K H2 layer (Herczeg et al. 2004; Hoadley et al. 2015); significant
additional broadening would require temperatures in excess of the $\approx$
4500 K dissociation temperature of H2 (Lepp & Shull 1983). If we further
assume that any turbulence in the disks is subsonic, then the maximum
turbulent velocity will be no larger than a few km s-1. Therefore, velocity
broadening due to bulk motions and Keplerian rotation dominate the observed
line shapes when the FWHM of the emission line is greater than the $\approx$
15 – 20 km s-1 spectral resolution of COS (see Figure 2; all H2 lines are
well-resolved at the spectral resolution of $HST$-COS; the 1450 Å $HST$+COS
LSF overplotted as the green dotted line). For the case of H2 in a
circumstellar disk where Keplerian broadening dominates the profile, we employ
a simple metric to characterize the average H2 emitting radius,
$\langle$$R_{H2}$$\rangle$ (Salyk et al. 2011a; France et al. 2012a;
Arulanantham et al. 2021), where

$\langle R_{H2}\rangle_{m}=GM_{*}\left(\frac{2sin(i)}{FWHM_{m}}\right)^{2}$
(4)

where $M_{*}$ is the stellar mass, $i$ is the disk inclination angle, taken
primarily from ALMA surveys (Table 1), and FWHMm is the mean of the Gaussian
FWHMs for a given progression $m$. When disk inclinations were not available
(19/63 accreting sources), we assume a disk inclination of 60°, as this is
approximately the average inclination found in ALMA disk surveys (Ansdell et
al. 2018) and assumed in model studies (Rilinger & Espaillat 2021). We provide
correlation coefficients for the full sample and using only the sources with
known inclinations in the appropriate subsections below, and note that the
“outlier” values of $\langle$$R_{H2}$$\rangle$ are not driven by unknown
inclinations; 5/6 of the largest H2 radii presented here are for sources with
measured disk inclinations. Table 1 presents the stellar masses and disk
inclinations used in this work.

For the H2 radial distribution analysis, we consider four progressions: the
[1,7] and [1,4] progressions that are pumped within 100 km s-1 of the
Ly$\alpha$ line center and two progressions ([0,1] and [0,2]) that are pumped
from the wing of the Ly$\alpha$ emission profile (379 and 487 km s-1 from
line-center, respectively). We analyze multiple progressions because self-
absorption can selectively impact the lower-energy level gas (Wood et al.
2002; McJunkin et al. 2016) and that off-star Ly$\alpha$ emission and the
local scattering geometry can impact the flux distributions of the different
progressions (Walter et al. 2003; Arulanantham et al. 2023). We will return to
a discussion of these results in Section 4.

###  3.3 H2 Dissociation Emission and the FUV Continuum

FUV continuum spectra were created for the ULLYSES DR5 sample following the
continuum extraction methodology described in France et al. (2014c). We create
a grid of 205 unique spectral points between 1138 and 1791 Å, selected by hand
to avoid discrete molecular and atomic emission and absorption features, where
0.75 Å (approximately 10 spectral resolution elements) spectral continuum
windows can be cleanly measured. We measure the mean and standard deviation of
the observed spectra, and these points define the binned flux spectrum and
error array (see Figure 3). These binned spectra are then corrected for
interstellar reddening and the FUV continuum and H2 continuum emission (the
“1600 Å Bump”) are separated as described below.

Figure 6: The FUV continuum luminosity increases with pre-main sequence
stellar mass. The stars in the ULLYSES sample with good measurements of the
average H2 emitting radius are shown with a black x, and those with literature
AV $<$ 1 are overlaid with pink circles (see Section 3.3 and 4.1). A fit to
the AV $<$ 1 points is overplotted in purple. The Spearman rank coefficient
($\rho$) and probability coefficient for null correlation (n) of the AV $<$ 1
points are shown at the upper left.

We measure the FUV continuum emission by fitting the binned spectrum with a
second order polynomial at wavelengths away from the Bump and residual stellar
or airglow emission, in the regions $\Delta$$\lambda$ = 1145 – 1190 Å, 1245 –
1330 Å, 1395 – 1401 Å, 1420 – 1465 Å, 1690 – 1710 Å, and 1730 – 1750 Å (see
the blue dashed lines in Figure 3). The polynomial fit is extrapolated to the
atomic hydrogen ionization limit, 912 Å, to cover the primary
photodissociating spectral region for H2, CO, and N2. Comparison with archival
$FUSE$ observations between 1000 and 1180 Å has shown that this extrapolation
is a reasonably good match to the observed short-wavelength flux (France et
al. 2014b). The FUV continuum flux, $F_{FUVCont}$, is the integral of the
polynomial over the 912 – 1760 Å wavelength region and the FUV continuum
luminosity is defined as $L$(FUV Cont) = 4$\pi$$d^{2}$$F_{FUVCont}$.

Figure 7: A comparison of the C IV luminosity with the average H2 emitting
radius for four of the progressions studied in this work (color-coded as
indicated in the legend). As has been shown in previous work, the majority of
H2 emission in CTTS systems arises between 0.1 and 3 au. A best fit power-law
relationship between L(C IV) and $\langle$$R_{H2}$$\rangle$ for the [1,4]
progression is shown overplotted in purple and the Spearman rank coefficient
($\rho$) and probability coefficient for null correlation (n) are shown at the
upper right. Assuming that C IV serves as a proxy for the accretion
luminosity, the H2 radial distribution demonstrates a weak dependence on the
strength of the high-energy emission, possibly suggesting increased molecular
dissociation in the innermost region of rapidly accreting systems.

Figure 6 shows the FUV continuum luminosity as a function of stellar mass. We
observe a positive correlation between the continuum and stellar mass ($\rho$
= 0.570, n = 1.0 $\times$ 10-4 and shown in the legend of Figure 6), which is
expected in the scenario where the FUV continuum emission is dominated by mass
accretion onto the central star (France et al. 2014c). Figure 6 plots the
sources for which “good” H2 radii are measured (see the discussion in Section
4.1), and it is observed that the most luminous sources are also the sources
with the largest reddening (AV). The attenuation towards FUV emitting regions
of protoplanetary targets has been estimated to be lower than estimates
inferred from X-ray, optical, or IR measurements (e.g., McJunkin et al. 2014).
Hence, extinction values derived from, e.g., IR observations may overestimate
the extinction correction for UV data. This was demonstrated by France et al.
(2017), who showed that the largest AV-values derived from IR-based
extinctions led to unphysical correlations in FUV luminosity relationships. To
avoid biasing the results by the largest, possibly incorrect, AV values,
Figure 6 also shows the 41 stars with AV $<$ 1 as pink circles. The
correlation coefficient and best-fit line are presented for the AV $<$ 1
sample.

The most prominent spectral feature in the binned spectra of some CTTSs is the
1600 Å Bump feature created during the dissociation of H2 molecules with
highly non-thermal rovibrational distributions, which spans $\sim$ 30 – 150 Å
and has a peak flux $\sim$ 0 – 4 times the underlying continuum level at 1600
Å. We define the 1600 Å Bump as the excess emission above the FUV continuum in
the 1490 – 1690 Å spectral region, where the Bump spectrum is the binned data
minus the FUV continuum fit. An example of the well-defined Bump spectrum of
DM Tau is shown in Figure 3. Following previous work, we parameterize the Bump
emission as a broad Gaussian plus a second-order polynomial (France et al.
2017, Figure 3, bottom panel). We define the total Bump flux, $F_{Bump}$, as
the integral of the continuum-subtracted Bump spectrum from 1490 to 1690 Å,
and $L$(1600 Å Bump) = 4$\pi$$d^{2}$$F_{Bump}$.

Figure 8: A comparison of the dust masses derived from mm-wave observations
(see Manara et al. 2022 for a review) with the average H2 emitting radius for
the [1,4] progression (29 sources had both dust masses, disk inclinations $>$
15°, and good H2 FWHM observations). The Spearman rank coefficient ($\rho$)
and probability coefficient for null correlation (n) are shown at the upper
left. The behavior is consistent with a picture where more massive stars
hosting more massive dust disks also photodissociate their inner few 0.1 au,
leading to a correlation between disk mass and H2 emitting radii.  Figure 9: A
comparison of the $WISE$ 12$\mu$m - 22$\mu$m color with the average H2
emitting radius for four of the progressions studied in this work (color-coded
as indicated in the legend). The distribution of emitting radii are clustered
between W3 - W4 $\approx$ 2.0 - 2.5 with a small number of transition disks
(Sz 111, DM Tau, CS Cha) extending to large MIR flux ratios. The Spearman rank
coefficient ($\rho$) and probability coefficient for null correlation (n) are
shown at the upper left.

##  4 Discussion: The Spatial Distribution and Excitation Conditions of H2
Line and Continuum Emission

###  4.1 Variation of $\langle$$R_{H2}$$\rangle$ with FUV Emission and Dust
Environments

The average H2 line emission radius represents the characteristic Keplerian
velocity of the gas, assuming the spectral broadening is dominated by the
rotational motion of the disk. In a future work, we will present a more
detailed line profile decomposition of a subset of bright disks to constrain
stratified gas structures and wind contributions (e.g., Herczeg et al. 2006;
Gangi et al. 2023), but for the full ULLYSES sample, we assume a single
emission line component. The objects included in the analysis are restricted
to disk inclinations $>$ 15 degrees (the Doppler motion cannot be confidently
measured for face-on disks at the resolution of $HST$-COS) and do not have any
data quality flags (low S/N, spectral blending) in the H2 emission line
catalog.

The final sample of sources that meet these criteria (“good” H2 radii)
includes 52, 38, 39, and 31 disks in the [1,4], [1,7], [0,1], and [0,2]
progressions, respectively. $\langle$$R_{H2}$$\rangle$ values range between
0.05 – 2.9 au, with average emitting radii between 0.53 – 0.72 au (median
radii between 0.35 – 0.57 au) for the four progressions. The majority of the
fluorescent line emission (based on the standard deviation of the radial
distributions of the four progressions) arises between 0.1 – 1.4 au. This
typical emitting range can be seen in the cluster of radial distributions in
Figure 7. The minimum range of H2 emitting radii (0.05 – 0.1 au) are
comparable to the dust sublimation radius and the inner disk dust rim radii
measured by GRAVITY (0.1 – 0.2 au; GRAVITY Collaboration et al. 2021),
suggesting that the molecular gas and dust edges may overlap in some sources.

Figure 7 shows the relationship between $\langle$$R_{H2}$$\rangle$ and the C
IV luminosity ($\rho$ = 0.383, n = 5.6 $\times$ 10-3, reducing to $\rho$ =
0.330, n = 4.9 $\times$ 10-2 if restricted to disks with measured
inclinations). Given the variability in mass accretion rates and FUV fluxes
over time (Espaillat et al. 2019; Hinton et al. 2022; Claes et al. 2022), we
use the $L$(C IV) measured contemporaneously with the H2 emission lines to
provide a proxy for the ’real-time’ FUV radiation field of the central star.
In Figure 7, we observe a marginal positive correlation between the emitting
radius and the strength of the C IV emission. This result may indicate that
higher FUV fluxes result in larger equilibrium radii with the
heating/dissociating FUV radiation field (e.g., Cazzoletti et al. 2018),
suggesting higher levels of H2 dissociation, or the clearing of H2-shielding
inner disk dust, close to the central star in more FUV-luminous systems
(Figure 6). We note that the lowest C IV luminosity systems (log10($L$(C IV))
$\lesssim$ 28.7) populate the lower left quadrant in Figure 7
(log10$\langle$$R_{H2}$$\rangle$ $\lesssim$ 0.3). Comparing the distribution
of $L$(C IV) with the stellar mass reveals that all stars with log10($L$(C
IV)) $\lesssim$ 28.7 also have M∗ $<$ 0.3 M⊙. This supports the idea that the
lowest mass sources with the lowest $L$(C IV) show less dissociation in their
innermost disks than their more massive counterparts. This would lead the
stellar Ly$\alpha$ photons to encounter optically thick gas at smaller radial
distances, which manifests as a small $\langle$$R_{H2}$$\rangle$ in these
systems.

We observe a stronger correlation ($\rho$ = 0.692, n = 3.2 $\times$ 10-5)
between the average H2 emitting radius and the disk dust mass (Figure 8). The
actual distribution is not a straight linear trend, instead, the H2 emitting
radii are roughly constant between $\approx$ 0.1 – 0.4 au for dust masses
$\lesssim$ 20 MEarth, with no H2 radii larger than $\approx$ 0.4 au. For dust
masses $\gtrsim$ 20 MEarth, the H2 radii range from 0.2 – 2 au. It is
established that more massive stars host massive dust disks (Andrews et al.
2013; Pascucci et al. 2016; Ansdell et al. 2017) and that dust disk mass is
correlated with mass accretion rate (Manara et al. 2016; Mulders et al. 2017;
Manara et al. 2022). The observation of larger gas clearing radii around stars
with larger dust masses (and higher FUV continuum luminosities) is consistent
with this picture of inner disk structure, as is the scenario presented above
where the stars with the lowest masses (and C IV luminosities) were found to
host H2 populations with the smallest emitting radii.

Figure 9 displays the relationship between the $WISE$ W3 – W4 color666WISE W3
is centered at 12 $\mu$m, full band range $\sim$ 7.5 – 16 $\mu$m; WISE W4 is
centered at 22 $\mu$m, full band range $\sim$ 20 – 25 $\mu$m; Wright et al.
(2010). and $\langle$$R_{H2}$$\rangle$. With the relatively low Spearman
correlation coefficient ($\rho$ = 0.326) and high probability of non-
correlation (n = 0.079), we do not find a correlation between the slope of the
inner disk dust SED and the average H2 emitting radius. This is in some
tension with the findings of Hoadley et al. (2015), who observed a positive
correlation between their modeled radial distribution of H2 fluorescence and
the $n_{13-31}$ dust SED derived from $Spitzer$-IRS spectroscopy. There are
two key differences between these studies: first, the $Spitzer$ data employed
by Hoadley et al. (2015) are a cleaner metric of warm inner disk dust clearing
than the $WISE$ photometric points (Espaillat et al. 2014). The $WISE$ W3 band
in particular includes the 10 $\mu$m silicate emission feature seen in many
CTTS spectra (Olofsson et al. 2010), and the lower angular resolution of
$WISE$ increases the possibility of source confusion. Second, Hoadley et al.
(2015) employ an emission line forward modeling analysis of a higher-S/N
sample of disk targets; they find correlations between inner disk dust
clearing and the maxima of the H2 radial distributions derived from their
models. A follow-on study of the high-S/N ULLYSES sources with $Spitzer$ or
$JWST$ spectra would be valuable for a more direct comparison with the Hoadley
et al. (2015) results.

###  4.2 Inner Disk Opacity and H2 Continuum Emission

There are two proposed origins for the 1600 Å Bump emission seen in many
protoplanetary disk sources. The first proposes that free electrons (possibly
liberated by X-ray emission from the central star) collide with the disk
surface population of H2, creating an electron impact spectrum similar to that
observed in the aurorae of giant planets (Bergin et al. 2004; Ingleby et al.
2009). The second route for the formation of the 1600 Å Bump proposes that
bright stellar+accretion-generated Ly$\alpha$ emission dissociates water
molecules and that the broad continuum Bump is produced during the
photodissociation of the non-thermal H2 created during the break-up
(dissociation) of the water molecules (van Harrevelt & van Hemert 2008; France
et al. 2017).

Analysis of the ULLYSES DR5 sample finds 24 sources with detectable 1600 Å
Bump emission, a detection rate of 38%, roughly consistent with the $\approx$
50% detection rate reported in France et al. (2017). We find a strong
correlation between the strength of the Bump and the Ly$\alpha$ field
strength. Figure 10 illustrates the relationship between the Ly$\alpha$ and
1600 Å Bump luminosity, the highly significant correlation ($\rho$ = 0.8, n =
2.7 $\times$ 10-6) strongly suggests Ly$\alpha$ photons as central to the
formation of the H2 continuum emission feature. For comparison, the
correlation between C IV and the 1600 Å Bump luminosity is less pronounced
($\rho$ = 0.5, n = 1.6 $\times$ 10-2). We also note that the Ly$\alpha$
luminosities in sources with strong Bump detections are $\gtrsim$ 5 $\times$
1031 erg s-1, the Ly$\alpha$ region strongly anchored by the France et al.
(2012a) Ly$\alpha$ reconstruction sample.

Figure 10: The 1600Å Bump luminosity is strongly correlated with Ly$\alpha$
luminosity. This correlation supports the hypothesis that the Bump is powered
by Ly$\alpha$-driven H2O dissociation. The Spearman rank coefficient ($\rho$)
and probability coefficient for null correlation (n) are shown at the upper
left.  Figure 11: A comparison between the $WISE$ 12$\mu$m - 22$\mu$m color
and the 1600Å Bump luminosity. Targets are labeled for reference and the
Spearman rank coefficient ($\rho$) and probability coefficient for null
correlation (n) are shown at the upper left. The positive correlation
indicates that as the inner disk opacity decreases, stellar+accretion-
generated Ly$\alpha$ photons can penetrate farther into the disk where they
dissociate H2O, creating the non-thermal H2 population that gives rise to the
1600 Å molecular dissociation continuum.

Within the sample of 24 1600 Å Bump detections, 14 had $WISE$ W3-W4 colors
without data quality warnings. We plot the Bump luminosity versus the W3-W4
color in Figure 11. We observe a correlation ($\rho$ = 0.68, n = 7.6 $\times$
10-3) between these quantities, with disks showing the largest depletion of
optically thick warm dust (noting the caveats about spectral and spatial
contamination in the $WISE$ photometric bands discussed previously) having the
largest H2 continuum emission luminosity. Taken together with the strong
correlation with L(Ly$\alpha$), the emerging picture for the formation of the
1600 Å Bump includes both strong accretion-generated Ly$\alpha$ flux and low
optical depth in the inner disk gas and dust environment. Bright Ly$\alpha$
emission alone is not a sufficient condition to produce the H2 continuum
emission, as there are 44 targets in our sample with L(Ly$\alpha$) $>$ 1031
erg s-1 (including, e.g., CVSO 109, Figure 3 and Espaillat et al. 2022) yet
only 21 of those are detected in the 1600 Å Bump feature. The second required
condition is that those Ly$\alpha$ photons are able to reach the molecular
reservoirs in the inner disk. As the gas and dust opacities decline with time,
Ly$\alpha$ photons can more readily reach the molecular populations of the
inner disk, where they dissociate water molecules and the resultant non-
thermal H2 gas.

##  5 Summary

We have analyzed far-ultraviolet spectra of 71 young stars included in the
ULLYSES Data Release 5. We have created a uniform emission line extraction and
measurement database of fluorescent H2 and magnetospheric and accretion-
generated hot gas lines that are hosted on a public website at the University
of Colorado3. The H2 database is used to compute the total fluorescent H2
luminosity and to estimate the Ly$\alpha$ luminosity from the ULLYSES sample.

The line-widths are combined with the stellar mass and disk inclination to
derive the average H2 emitting radius for the fluorescent emission. We find
that the majority of the fluorescent emission arises between 0.1 and 1.4 au
from the central star. We find a weak positive correlation between the average
emitting radius and the stellar FUV flux, and a stronger correlation between
the H2 emitting radius and the dust disk mass. We extracted FUV continuum
spectra and H2 dissociation emission from the ULLYSES sample, finding 24
detections of H2 continuum emission. We find strong correlations between the
H2 continuum emission and both the Ly$\alpha$ luminosity and the mid-IR (10 –
22 $\mu$m) inner disk dust slope. These results support a picture for the
production of the H2 continuum where a decrease of the inner disk gas and dust
opacity enables bright stellar Ly$\alpha$ emission to reach the molecular
populations of the inner disk (Bergin et al. 2003; Bethell & Bergin 2011),
driving molecular dissociation and the associated continuum emission.

This work was carried out as part of the Outflows and Disks around Young
Stars: Synergies for the Exploration of Ullyses Spectra (ODYSSEUS) program,
supported by $HST$ archival grant number 16129-025 to the University of
Colorado at Boulder. E.F. acknowledges support from the project PRIN‐INAF 2019
“Spectroscopically Tracing the Disk Dispersal Evolution (STRADE)”. We
acknowledge the excellent work of the ULLYSES implementation team at STScI to
develop and execute this ambitious spectroscopic observing program. The HST-
ULLYSES data presented in this paper were obtained from the Mikulski Archive
for Space Telescopes (MAST) at the Space Telescope Science Institute. The
specific observations analyzed can be accessed via
https://doi.org/10.17909/t9-jzeh-xy14 (catalog DOI: 10.17909/t9-jzeh-xy14).

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Table 1: ULLYSES DR5 Target Parameters

Name | M∗ | $\dot{M}_{acc}$ | $i$ | $d$ | AV |  $WISE$ W3 |  $WISE$ W4 | References
---|---|---|---|---|---|---|---|---
| (M⊙) | (10-8 M⊙ yr-1) | ($deg$) | (pc) | (mag) | (mag) | (mag) |
2MASS-j04390163+2336029 | 0.17 | 0.04 | $\cdots$ | 126 | 0.00 | $\cdots$ | $\cdots$ | 14,15
2MASSj11432669-7804454 | 0.14 | 0.19 | $\cdots$ | 153 | 0.40 | 8.676 | 7.052 | 1
CHX18n | 0.81 | 0.47 | 56 | 192 | 0.80 | 5.560 | 3.584 | 1,29
CVSO-17 | 0.37 | 0.00 | $\cdots$ | 414 | 0.00 | 11.130 | 8.932 | 7
CVSO-36 | 0.39 | 0.00 | $\cdots$ | 335 | 0.10 | 10.306 | 8.593 | 7
CVSO-58 | 0.81 | 0.43 | $\cdots$ | 349 | 0.80 | 7.000 | 4.905 | 7
CVSO-90 | 0.62 | 0.25 | $\cdots$ | 338 | 0.10 | 7.622 | 5.345 | 7
CVSO-104 | 0.37 | 0.32 | 43 | 360 | 0.20 | 7.560 | 5.571 | 7,20
CVSO-107 | 0.53 | 5.01 | $\cdots$ | 330 | 0.30 | 7.419 | 5.128 | 7
CVSO-109 | 0.46 | 3.24 | $\cdots$ | 400 | 0.10 | 6.523 | 4.527 | 7
CVSO-146 | 0.86 | 0.27 | $\cdots$ | 332 | 0.60 | 7.072 | 4.597 | 7
CVSO-165 | 0.84 | 0.08 | $\cdots$ | 400 | 0.20 | 6.571 | 4.626 | 7
CVSO-176 | 0.25 | 1.45 | $\cdots$ | 302 | 1.00 | 7.896 | 6.221 | 7
echa-j0843-7915 (RECX-15) | 0.20 | 0.08 | 60 | 103 | 0.00 | $\cdots$ | $\cdots$ | 4,5,27
echa-j0844-7833 | 0.07 | 0.00 | $\cdots$ | 98 | 0.50 | $\cdots$ | $\cdots$ | 4,5
LkCa-4 | 0.77 | 0.00 | $\cdots$ | 129 | 0.69 | 8.038 | 7.934 | 9,10
LkCa-15 | 0.85 | 0.13 | 49 | 157 | 0.60 | 5.696 | 3.565 | 11,3,19
LkCa-19 | 1.35 | 0.00 | $\cdots$ | 157 | 0.00 | 7.972 | 7.109 | 9
HN5 | 0.18 | 0.28 | $\cdots$ | 195 | 1.10 | 5.864 | 4.371 | 1
RECX-1 | 0.75 | 0.00 | $\cdots$ | 98 | 0.00 | 7.058 | 6.943 | 4,5
RECX-11 | 0.83 | 0.02 | 70 | 99 | 0.10 | 5.388 | 3.718 | 4,5,27
RXJ0438.6+1556 | 1.20 | 0.03 | $\cdots$ | 139 | 0.30 | 8.154 | 8.122 | 13
RXJ1556.1-3655 | 0.50 | 1.20 | 53 | 158 | 1.00 | 6.268 | 3.887 | 16,30
SSTC2dj160000-422158 | 0.19 | 0.06 | 65 | 159 | 0.10 | $\cdots$ | $\cdots$ | 16,29
SSTC2dj160830-382827 | 1.40 | 0.11 | 72 | 153 | 0.20 | $\cdots$ | $\cdots$ | 16, 17
SSTC2dj161243-381503 | 0.44 | 0.17 | 43 | 159 | 0.80 | $\cdots$ | $\cdots$ | 16,30
SSTC2dj161344-373646 | 0.16 | 0.11 | 54 | 158 | 0.60 | $\cdots$ | $\cdots$ | 16,29
Sz 19 | 2.08 | 2.34 | 45 | 189 | 1.50 | 3.011 | 1.057 | 1,29
Sz 45 | 0.56 | 0.51 | 43 | 189 | 0.70 | 6.948 | 4.156 | 1,29
Sz 66 | 0.29 | 0.05 | 40 | 156 | 1.00 | 6.363 | 4.432 | 6,29
Sz 68 | 1.40 | 0.58 | 32 | 152 | 1.00 | 2.896 | 0.856 | 16,30
Sz 69 | 0.20 | 0.03 | 35 | 155 | 0.00 | 5.614 | 4.420 | 6,28
Sz 71 | 0.37 | 0.11 | 40 | 156 | 0.70 | 5.716 | 3.671 | 6,28
Sz 72 | 0.37 | 0.95 | 53 | 156 | 1.00 | 5.987 | 3.911 | 6,28
Sz 75 | 0.82 | 6.79 | $\cdots$ | 152 | 1.00 | 4.342 | 2.379 | 16
Sz 76 | 0.22 | 0.07 | $\cdots$ | 159 | 0.30 | 7.366 | 5.274 | 16
Sz 77 | 0.75 | 0.07 | $\cdots$ | 155 | 0.30 | 5.362 | 3.437 | 16
Sz 84 | 0.16 | 0.06 | 75 | 155 | 0.00 | 9.115 | 6.340 | 6,29
Sz 97 | 0.23 | 0.03 | 73 | 157 | 0.00 | 7.749 | 4.976 | 6,29
Sz 98 | 0.70 | 5.89 | 46 | 156 | 1.00 | 3.986 | 2.097 | 16,29
Sz 99 | 0.23 | 0.04 | $\cdots$ | 158 | 0.00 | 7.847 | 6.009 | 6
Sz 100 | 0.16 | 0.04 | 47 | 141 | 0.00 | 6.486 | 4.517 | 6,29
Sz 102 | 0.24 | 0.08 | 53 | 160 | 1.13 | 6.196 | 3.382 | 16,29
Sz 103 | 0.22 | 0.10 | 50 | 157 | 0.70 | 6.940 | 4.843 | 6,29
Sz 104 | 0.16 | 0.02 | 57 | 159 | 0.00 | 7.438 | 4.806 | 6,29
Sz 110 | 0.22 | 0.30 | 32 | 157 | 0.00 | 6.307 | 4.116 | 6,29
Sz 111 | 0.58 | 0.09 | 54 | 158 | 0.50 | 8.922 | 5.731 | 6,29
Sz 114 | 0.21 | 0.11 | 6 | 153 | 0.30 | 5.721 | 3.380 | 6,29
Sz 117 | 0.26 | 0.25 | 47 | 156 | 0.50 | 7.169 | 4.884 | 16,29
Sz 129 | 0.79 | 0.40 | 31 | 160 | 0.90 | 5.722 | 3.597 | 16,29
Sz 130 | 0.36 | 0.15 | 36 | 159 | 0.40 | 6.530 | 4.514 | 6,28
AA Tau | 0.80 | 0.33 | 59 | 134 | 0.50 | 4.645 | 2.505 | 11,3,18
CE Ant | 0.46 | 0.00 | $\cdots$ | 34 | 0.00 | 6.598 | 5.998 | 3
CS Cha | 1.05 | 1.20 | 37 | 168 | 0.80 | 7.095 | 2.758 | 2,3,13,19
DE Tau | 0.59 | 2.64 | 34 | 128 | 0.60 | 4.895 | 2.725 | 11,3,21
DK Tau | 0.71 | 3.79 | 58 | 140 | 0.80 | 3.599 | 1.800 | 11,3,22
DM Tau | 0.50 | 0.29 | 35 | 144 | 0.00 | 7.085 | 3.571 | 2,3,24
DN Tau | 0.60 | 0.35 | 35 | 128 | 1.90 | 5.166 | 3.039 | 11,3,23
DR Tau | 0.80 | 3.16 | 5 | 193 | 0.48 | 3.008 | 1.066 | 11,3,30
HN Tau | 0.85 | 0.13 | 50 | 140 | 0.50 | 4.174 | 2.160 | 11,3,31
IN Cha | 0.19 | 0.05 | $\cdots$ | 193 | 0.20 | 7.218 | 5.316 | 1
IP Tau | 0.68 | 0.08 | 35 | 129 | 0.20 | 5.436 | 3.578 | 11,3,31
MY Lup | 1.06 | 0.02 | 72 | 157 | 1.30 | 5.164 | 2.833 | 16,26
RY Lup | 1.71 | 1.00 | 68 | 153 | 0.40 | 3.647 | 1.404 | 16,19
SY Cha | 0.78 | 0.04 | 51 | 180 | 0.50 | 5.435 | 3.212 | 1,29
TX Ori | 1.09 | 6.17 | $\cdots$ | 385 | 0.40 | 5.101 | 3.083 | 8
UX Tau A | 1.30 | 1.00 | 35 | 140 | 0.20 | 5.710 | 2.176 | 12,19
V397 Aur | 0.70 | 0.00 | $\cdots$ | 165 | 0.20 | 7.854 | 6.743 | 9
XX Cha | 0.29 | 0.10 | $\cdots$ | 192 | 0.30 | 5.751 | 3.569 | 1
V505 Ori | 0.81 | 0.30 | $\cdots$ | 392 | 1.00 | 6.564 | 4.539 | 8
V510 Ori | 0.76 | 0.58 | $\cdots$ | 390 | 0.10 | 5.598 | 3.258 | 8

Note. — References for the stellar masses, reddening, accretion rates, and
disk inclination: (1) (Manara et al. 2017b), (2) (Furlan et al. 2009), (3)
(Ingleby et al. 2013), (4) (Rugel et al. 2018), (5) (Luhman 2004), (6) (Alcalá
et al. 2014), (7) (Calvet et al. 2005), (8) (Maucó et al. 2016), (9) (Kenyon &
Hartmann 1995), (10) (Donati et al. 2014) (11) (Furlan et al. 2011) (12)
(Costigan et al. 2014), (13) (Manara et al. 2017a), (14) (Slesnick et al.
2006), (15) (Herczeg & Hillenbrand 2008), (16) (Alcalá et al. 2017), (17) (van
der Marel et al. 2018), (18) (Loomis et al. 2017), (19) (Francis & van der
Marel 2020), (20) (Frasca et al. 2021), (21) (Simon et al. 2019), (22)
(Nelissen et al. 2023), (23) (Long et al. 2019), (24) (Kudo et al. 2018), (25)
(Kurtovic et al. 2018), (26) (Huang et al. 2018), (27) (Lawson et al. 2004),
(28) (Ansdell et al. 2016), (29) (Hendler et al. 2020), (30) (Braun et al.
2021), (31) (Simon et al. 2017)

Table 2: ULLYSES FUV Measurements Name | L(H2) | L(Ly$\alpha$) | L(C IV) | L(1600Å Bump) | L(FUV Cont) | FWHM(H2)[1,4] | FWHM(H2)[0,1] | $<R(H_{2})>_{[1,4]}$ | $<R(H_{2})>_{[0,1]}$
---|---|---|---|---|---|---|---|---|---
| (1029 erg s-1) | (1029 erg s-1) | (1029 erg s-1) | (1029 erg s-1) | (1029 erg s-1) | (km s-1) | (km s-1) | (au) | (au)
2MASS-j04390163+2336029 | 0.05 $\pm$ 0.01 | 2.89 | 0.02 |  $\cdots$ $\pm$ $\cdots$ | 0.1 | 76.4 $\pm$ 17.3 | 13.4 $\pm$ 29.0 | 0.08 $\pm$ 0.02 | 2.52 $\pm$ 5.46
2MASSj11432669-7804454 | 2.20 $\pm$ 0.22 | 61.18 | 0.65 | 0.49 $\pm$ 0.28 | 3.6 | 61.7 $\pm$ 14.8 | 67.0 $\pm$ 4.5 | 0.10 $\pm$ 0.02 | 0.08 $\pm$ 0.01
CHX18n | 17.65 $\pm$ 2.75 | 348.59 | 8.38 |  $\cdots$ $\pm$ $\cdots$ | 36.8 | 51.8 $\pm$ 4.9 | 49.1 $\pm$ 4.5 | 0.75 $\pm$ 0.07 | 0.83 $\pm$ 0.08
CVSO-17 | 0.04 $\pm$ 0.07 | 15.01 | 0.17 |  $\cdots$ $\pm$ $\cdots$ | 4.2 |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$
CVSO-36 | 0.11 $\pm$ 0.09 | 11.62 | 0.09 |  $\cdots$ $\pm$ $\cdots$ | 2.8 |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$
CVSO-58 | 37.82 $\pm$ 3.82 | 817.71 | 21.69 | 10.30 $\pm$ 10.00 | 277.6 | 54.3 $\pm$ 2.5 | 44.7 $\pm$ 4.4 | 0.73 $\pm$ 0.03 | 1.08 $\pm$ 0.11
CVSO-90 | 16.74 $\pm$ 1.74 | 423.96 | 10.30 |  $\cdots$ $\pm$ $\cdots$ | 79.1 | 51.8 $\pm$ 4.9 | 51.7 $\pm$ 1.6 | 0.62 $\pm$ 0.06 | 0.62 $\pm$ 0.02
CVSO-104 | 8.66 $\pm$ 0.69 | 258.39 | 4.14 |  $\cdots$ $\pm$ $\cdots$ | 51.8 | 49.3 $\pm$ 2.5 | 53.7 $\pm$ 4.2 | 0.25 $\pm$ 0.01 | 0.21 $\pm$ 0.02
CVSO-107 | 15.21 $\pm$ 2.49 | 388.81 | 7.94 |  $\cdots$ $\pm$ $\cdots$ | 136.0 | 64.1 $\pm$ 7.4 | 53.6 $\pm$ 13.4 | 0.34 $\pm$ 0.04 | 0.49 $\pm$ 0.12
CVSO-109 | 14.60 $\pm$ 1.52 | 407.72 | 12.72 |  $\cdots$ $\pm$ $\cdots$ | 229.5 | 44.4 $\pm$ 2.5 | 40.7 $\pm$ 3.0 | 0.62 $\pm$ 0.03 | 0.74 $\pm$ 0.05
CVSO-146 | 28.78 $\pm$ 4.50 | 645.24 | 15.46 |  $\cdots$ $\pm$ $\cdots$ | 310.0 | 51.8 $\pm$ 19.7 | 53.6 $\pm$ 6.7 | 0.85 $\pm$ 0.32 | 0.80 $\pm$ 0.10
CVSO-165 | 11.39 $\pm$ 2.00 | 335.04 | 5.63 |  $\cdots$ $\pm$ $\cdots$ | 66.2 | 76.4 $\pm$ 2.5 | 53.6 $\pm$ 13.4 | 0.38 $\pm$ 0.01 | 0.78 $\pm$ 0.19
CVSO-176 | 51.99 $\pm$ 5.58 | 990.56 | 10.54 |  $\cdots$ $\pm$ $\cdots$ | 245.9 | 74.0 $\pm$ 2.5 | 46.9 $\pm$ 2.2 | 0.12 $\pm$ 0.00 | 0.30 $\pm$ 0.01
echa-j0843-7915 | 5.33 $\pm$ 0.48 | 104.36 | 0.48 | 0.96 $\pm$ 0.15 | 5.5 | 49.3 $\pm$ 2.5 | 55.8 $\pm$ 2.2 | 0.22 $\pm$ 0.01 | 0.17 $\pm$ 0.01
echa-j0844-7833’ | 0.15 $\pm$ 0.06 | 6.21 | 0.21 |  $\cdots$ $\pm$ $\cdots$ | 0.7 |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$
LkCa-4 | 0.07 $\pm$ 0.30 | 3.45 | 1.50 |  $\cdots$ $\pm$ $\cdots$ | 2.3 |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$
LkCa-15 | 23.78 $\pm$ 3.11 | 405.89 | 12.40 |  $\cdots$ $\pm$ $\cdots$ | 50.5 | 51.8 $\pm$ 2.5 | 53.6 $\pm$ 2.2 | 0.64 $\pm$ 0.03 | 0.60 $\pm$ 0.02
LkCa-19 | 0.08 $\pm$ 0.08 | 4.31 | 0.71 |  $\cdots$ $\pm$ $\cdots$ | 1.8 |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$
HN5 | 2.52 $\pm$ 0.35 | 75.19 | 0.41 |  $\cdots$ $\pm$ $\cdots$ | 5.4 | 133.2 $\pm$ 93.7 | 17.9 $\pm$ 2.2 | 0.03 $\pm$ 0.02 | 1.50 $\pm$ 0.19
RECX-1 | -0.10 $\pm$ 0.13 | 0.07 | 0.51 |  $\cdots$ $\pm$ $\cdots$ | 2.3 |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$
RECX-11 | 2.07 $\pm$ 0.26 | 48.42 | 2.35 | 0.44 $\pm$ 0.16 | 5.1 | 61.6 $\pm$ 2.5 | 65.9 $\pm$ 2.0 | 0.69 $\pm$ 0.03 | 0.60 $\pm$ 0.02
RXJ0438.6+1556 | 0.06 $\pm$ 0.10 | 3.10 | 1.01 |  $\cdots$ $\pm$ $\cdots$ | 2.7 | 182.5 $\pm$ 88.8 | 78.2 $\pm$ 67.0 | 0.10 $\pm$ 0.05 | 0.52 $\pm$ 0.45
RXJ1556.1-3655 | 105.72 $\pm$ 9.49 | 1323.61 | 53.50 |  $\cdots$ $\pm$ $\cdots$ | 116.5 | 56.6 $\pm$ 2.5 | 49.1 $\pm$ 2.2 | 0.36 $\pm$ 0.02 | 0.48 $\pm$ 0.02
SSTC2dj160000-422158 | 0.14 $\pm$ 0.01 | 6.88 | 0.02 |  $\cdots$ $\pm$ $\cdots$ | 0.5 | 61.6 $\pm$ 2.5 | 35.7 $\pm$ 20.1 | 0.15 $\pm$ 0.01 | 0.44 $\pm$ 0.25
SSTC2dj160830-382827 | 4.83 $\pm$ 0.75 | 113.89 | 1.28 |  $\cdots$ $\pm$ $\cdots$ | 6.6 | 46.8 $\pm$ 2.5 | 40.2 $\pm$ 20.1 | 2.06 $\pm$ 0.11 | 2.79 $\pm$ 1.39
SSTC2dj161243-381503 | 7.63 $\pm$ 0.46 | 166.38 | 3.83 | 0.82 $\pm$ 0.52 | 14.0 | 54.3 $\pm$ 4.9 | 42.4 $\pm$ 2.2 | 0.25 $\pm$ 0.02 | 0.41 $\pm$ 0.02
SSTC2dj161344-373646 | 2.17 $\pm$ 0.51 | 61.26 | 4.84 |  $\cdots$ $\pm$ $\cdots$ | 13.2 | 49.3 $\pm$ 2.5 | 35.7 $\pm$ 13.4 | 0.16 $\pm$ 0.01 | 0.30 $\pm$ 0.11
Sz 19 | 8.29 $\pm$ 1.69 | 190.54 | 13.98 |  $\cdots$ $\pm$ $\cdots$ | 42.1 | 49.3 $\pm$ 9.9 | 41.3 $\pm$ 4.5 | 1.44 $\pm$ 0.29 | 2.05 $\pm$ 0.22
Sz 45 | 21.74 $\pm$ 1.96 | 408.30 | 8.43 |  $\cdots$ $\pm$ $\cdots$ | 66.8 | 46.8 $\pm$ 2.5 | 42.4 $\pm$ 2.2 | 0.43 $\pm$ 0.02 | 0.53 $\pm$ 0.03
Sz 66 | 17.24 $\pm$ 1.59 | 313.71 | 0.11 | 2.64 $\pm$ 1.45 | 11.9 | 44.4 $\pm$ 7.4 | 44.7 $\pm$ 2.2 | 0.22 $\pm$ 0.04 | 0.22 $\pm$ 0.01
Sz 68 | 15.88 $\pm$ 2.39 | 291.37 | 8.66 | 3.74 $\pm$ 2.37 | 30.3 | 59.2 $\pm$ 12.3 | 58.1 $\pm$ 6.7 | 0.42 $\pm$ 0.09 | 0.43 $\pm$ 0.05
Sz 69 | 1.69 $\pm$ 0.21 | 49.80 | 0.03 | 0.50 $\pm$ 0.21 | 1.1 | 54.3 $\pm$ 4.9 | 42.4 $\pm$ 2.2 | 0.08 $\pm$ 0.01 | 0.13 $\pm$ 0.01
Sz 71 | 6.71 $\pm$ 0.74 | 148.76 | 2.58 | 1.25 $\pm$ 0.78 | 10.8 | 46.9 $\pm$ 2.5 | 38.0 $\pm$ 4.4 | 0.25 $\pm$ 0.01 | 0.38 $\pm$ 0.04
Sz 72 | 11.27 $\pm$ 1.98 | 224.08 | 14.26 |  $\cdots$ $\pm$ $\cdots$ | 65.0 | 56.7 $\pm$ 9.9 | 45.4 $\pm$ 1.9 | 0.26 $\pm$ 0.05 | 0.41 $\pm$ 0.02
Sz 75 | 162.14 $\pm$ 25.70 | 1826.48 | 73.94 |  $\cdots$ $\pm$ $\cdots$ | 966.3 | 78.9 $\pm$ 4.9 | 60.3 $\pm$ 6.7 | 0.35 $\pm$ 0.02 | 0.61 $\pm$ 0.07
Sz 76 | 1.33 $\pm$ 0.14 | 41.62 | 0.27 |  $\cdots$ $\pm$ $\cdots$ | 2.3 | 37.0 $\pm$ 2.5 | 60.3 $\pm$ 13.4 | 0.43 $\pm$ 0.03 | 0.16 $\pm$ 0.04
Sz 77 | 6.27 $\pm$ 0.78 | 140.53 | 2.77 | 1.09 $\pm$ 0.92 | 16.7 | 54.3 $\pm$ 9.9 | 46.9 $\pm$ 64.8 | 0.68 $\pm$ 0.12 | 0.91 $\pm$ 1.25
Sz 84 | 0.62 $\pm$ 0.09 | 22.58 | 0.10 |  $\cdots$ $\pm$ $\cdots$ | 0.6 | 44.4 $\pm$ 4.9 | 46.9 $\pm$ 11.1 | 0.27 $\pm$ 0.03 | 0.24 $\pm$ 0.06
Sz 97 | 0.29 $\pm$ 0.05 | 12.25 | 0.26 |  $\cdots$ $\pm$ $\cdots$ | 1.6 | 51.8 $\pm$ 2.5 | 105.0 $\pm$ 22.4 | 0.28 $\pm$ 0.01 | 0.07 $\pm$ 0.01
Sz 98 | 28.07 $\pm$ 2.34 | 461.61 | 4.79 | 2.32 $\pm$ 2.15 | 59.7 | 49.3 $\pm$ 2.5 | 44.7 $\pm$ 2.2 | 0.54 $\pm$ 0.03 | 0.66 $\pm$ 0.03
Sz 99 | 0.25 $\pm$ 0.05 | 10.95 | 0.41 |  $\cdots$ $\pm$ $\cdots$ | 1.7 | 61.5 $\pm$ 17.2 | 22.3 $\pm$ 58.0 | 0.16 $\pm$ 0.05 | 1.23 $\pm$ 3.19
Sz 100 | 1.30 $\pm$ 0.12 | 39.04 | 0.17 |  $\cdots$ $\pm$ $\cdots$ | 1.6 | 46.8 $\pm$ 2.5 | 15.8 $\pm$ 2.1 | 0.14 $\pm$ 0.01 | 1.22 $\pm$ 0.17
Sz 102 | 111.34 $\pm$ 12.58 | 1386.21 | 44.74 | 10.10 $\pm$ 6.69 | 379.8 | 59.1 $\pm$ 4.9 | 62.5 $\pm$ 2.2 | 0.16 $\pm$ 0.01 | 0.14 $\pm$ 0.00
Sz 103 | 7.63 $\pm$ 0.66 | 165.16 | 0.22 |  $\cdots$ $\pm$ $\cdots$ | 27.0 | 49.3 $\pm$ 2.5 | 53.6 $\pm$ 2.2 | 0.19 $\pm$ 0.01 | 0.16 $\pm$ 0.01
Sz 104 | 0.14 $\pm$ 0.07 | 6.95 | 0.21 |  $\cdots$ $\pm$ $\cdots$ | 3.5 | 76.3 $\pm$ 27.1 | 136.2 $\pm$ 109.4 | 0.07 $\pm$ 0.02 | 0.02 $\pm$ 0.02
Sz 110 | 2.00 $\pm$ 0.35 | 57.43 | 0.72 |  $\cdots$ $\pm$ $\cdots$ | 2.5 | 68.9 $\pm$ 9.9 | 100.5 $\pm$ 29.0 | 0.05 $\pm$ 0.01 | 0.02 $\pm$ 0.01
Sz 111 | 13.88 $\pm$ 1.54 | 265.72 | 3.30 | 4.05 $\pm$ 2.31 | 16.6 | 46.8 $\pm$ 2.5 | 58.1 $\pm$ 6.7 | 0.63 $\pm$ 0.03 | 0.41 $\pm$ 0.05
Sz 114 | 12.68 $\pm$ 0.97 | 244.59 | 1.51 |  $\cdots$ $\pm$ $\cdots$ | 5.6 | 44.3 $\pm$ 2.5 | 35.7 $\pm$ 2.2 | 0.01 $\pm$ 0.00 | 0.01 $\pm$ 0.00
Sz 117 | 1.36 $\pm$ 0.82 | 42.13 | 2.54 |  $\cdots$ $\pm$ $\cdots$ | 8.1 | 73.8 $\pm$ 152.6 | 78.1 $\pm$ 73.7 | 0.09 $\pm$ 0.19 | 0.08 $\pm$ 0.08
Sz 129 | 37.69 $\pm$ 4.56 | 588.70 | 18.31 | 8.39 $\pm$ 1.74 | 190.6 | 51.8 $\pm$ 2.5 | 49.1 $\pm$ 2.2 | 0.28 $\pm$ 0.01 | 0.31 $\pm$ 0.01
Sz 130 | 2.28 $\pm$ 0.35 | 63.87 | 1.33 |  $\cdots$ $\pm$ $\cdots$ | 5.0 | 49.3 $\pm$ 2.5 | 42.4 $\pm$ 22.3 | 0.19 $\pm$ 0.01 | 0.25 $\pm$ 0.13
AA Tau | 25.65 $\pm$ 2.60 | 403.91 | 2.04 | 2.70 $\pm$ 0.41 | 16.6 | 54.2 $\pm$ 2.5 | 53.6 $\pm$ 2.2 | 0.71 $\pm$ 0.03 | 0.73 $\pm$ 0.03
CE Ant | 0.00 $\pm$ 0.01 | 0.13 | 0.06 |  $\cdots$ $\pm$ $\cdots$ | 0.3 |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$
CS Cha | 133.10 $\pm$ 12.53 | 1632.53 | 0.44 | 27.60 $\pm$ 9.81 | 330.3 | 35.6 $\pm$ 0.3 | 35.7 $\pm$ 2.2 | 1.06 $\pm$ 0.01 | 1.06 $\pm$ 0.07
DE Tau | 12.17 $\pm$ 2.51 | 219.19 | 6.96 | 1.58 $\pm$ 0.64 | 40.5 | 59.2 $\pm$ 4.9 | 38.0 $\pm$ 8.9 | 0.19 $\pm$ 0.02 | 0.46 $\pm$ 0.11
DK Tau | 23.13 $\pm$ 2.96 | 378.29 | 5.64 | 1.33 $\pm$ 1.00 | 63.8 | 69.0 $\pm$ 7.4 | 35.7 $\pm$ 2.2 | 0.03 $\pm$ 0.00 | 0.10 $\pm$ 0.01
DM Tau | 5.54 $\pm$ 0.71 | 123.69 | 1.68 | 7.82 $\pm$ 2.41 | 13.6 | 44.5 $\pm$ 0.5 | 40.2 $\pm$ 4.5 | 0.29 $\pm$ 0.00 | 0.36 $\pm$ 0.04
DN Tau | 101.03 $\pm$ 26.50 | 1171.54 | 259.67 |  $\cdots$ $\pm$ $\cdots$ | 2580.2 | 41.9 $\pm$ 7.4 | 49.7 $\pm$ 3.1 | 0.40 $\pm$ 0.07 | 0.29 $\pm$ 0.02
DR Tau | 21.79 $\pm$ 7.61 | 412.74 | 11.40 |  $\cdots$ $\pm$ $\cdots$ | 2.9 $\times$ 105 | 39.4 $\pm$ 2.5 | 26.8 $\pm$ 13.4 | 0.02 $\pm$ 0.00 | 0.04 $\pm$ 0.02
HN Tau | 15.87 $\pm$ 2.66 | 280.88 | 2.98 | 2.11 $\pm$ 0.73 | 50.3 | 64.1 $\pm$ 4.9 | 53.6 $\pm$ 6.7 | 0.43 $\pm$ 0.03 | 0.62 $\pm$ 0.08
IN Cha | 0.71 $\pm$ 0.08 | 27.38 | 0.04 |  $\cdots$ $\pm$ $\cdots$ | 0.8 | 61.7 $\pm$ 2.5 | 40.2 $\pm$ 15.6 | 0.13 $\pm$ 0.01 | 0.31 $\pm$ 0.12
IP Tau | 1.08 $\pm$ 0.16 | 32.38 | 1.07 |  $\cdots$ $\pm$ $\cdots$ | 4.0 | 51.7 $\pm$ 4.9 | 53.6 $\pm$ 35.7 | 0.30 $\pm$ 0.03 | 0.28 $\pm$ 0.18
MY Lup | 68.89 $\pm$ 5.93 | 941.26 | 19.46 |  $\cdots$ $\pm$ $\cdots$ | 111.8 | 39.5 $\pm$ 2.5 | 40.2 $\pm$ 4.5 | 2.20 $\pm$ 0.14 | 2.12 $\pm$ 0.24
RY Lup | 10.90 $\pm$ 2.18 | 216.77 | 4.02 | 3.33 $\pm$ 2.01 | 18.7 | 49.3 $\pm$ 2.5 | 42.4 $\pm$ 13.4 | 2.15 $\pm$ 0.11 | 2.90 $\pm$ 0.92
SY Cha | 12.84 $\pm$ 1.15 | 264.34 | 2.87 | 6.01 $\pm$ 0.81 | 44.5 | 74.0 $\pm$ 2.5 | 55.8 $\pm$ 4.5 | 0.31 $\pm$ 0.01 | 0.55 $\pm$ 0.04
TX Ori | 11.38 $\pm$ 1.46 | 329.57 | 1.59 |  $\cdots$ $\pm$ $\cdots$ | 263.0 | 56.7 $\pm$ 2.5 | 40.2 $\pm$ 4.5 | 0.90 $\pm$ 0.04 | 1.80 $\pm$ 0.20
UX Tau A | 8.11 $\pm$ 0.59 | 165.18 | 2.29 | 1.10 $\pm$ 0.58 | 8.8 | 44.4 $\pm$ 2.5 | 44.7 $\pm$ 2.2 | 0.77 $\pm$ 0.04 | 0.76 $\pm$ 0.04
V397 Aur | -0.00 $\pm$ 0.05 | 2.22 | 0.45 |  $\cdots$ $\pm$ $\cdots$ | 0.4 |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$ |  $\cdots$ $\pm$ $\cdots$
XX Cha | 8.63 $\pm$ 1.09 | 198.02 | 2.72 |  $\cdots$ $\pm$ $\cdots$ | 27.6 | 51.8 $\pm$ 7.4 | 40.2 $\pm$ 2.2 | 0.29 $\pm$ 0.04 | 0.48 $\pm$ 0.03
V505 Ori | 49.55 $\pm$ 4.61 | 1063.19 | 2.93 |  $\cdots$ $\pm$ $\cdots$ | 1164.4 | 51.8 $\pm$ 4.9 | 38.0 $\pm$ 13.4 | 0.80 $\pm$ 0.08 | 1.50 $\pm$ 0.53
V510 Ori | 25.63 $\pm$ 4.81 | 629.94 | 1.62 | 15.50 $\pm$ 5.85 | 174.4 | 41.9 $\pm$ 4.9 | 42.4 $\pm$ 6.7 | 1.15 $\pm$ 0.13 | 1.12 $\pm$ 0.18