Berg et al.
CHAOS viii: Far-Ultraviolet Spectra of M101 and The Impact of Wolf-Rayet Stars[
Based on observations made with the NASA/ESA Hubble Space Telescope,
obtained from the Data Archive at the Space Telescope Science Institute, which
is operated by the Association of Universities for Research in Astronomy, Inc.,
under NASA contract NAS 5-26555.]

0000-0002-4153-053X]Danielle A. Berg
Department of Astronomy, The University of Texas at Austin, 2515 Speedway, Stop C1400, Austin, TX 78712, USA

0000-0003-0605-8732]Evan D. Skillman
Minnesota Institute for Astrophysics, University of Minnesota, 116 Church St. SE, Minneapolis, MN 55455

0000-0002-0302-2577]John Chisholm
Department of Astronomy, The University of Texas at Austin, 2515 Speedway, Stop C1400, Austin, TX 78712, USA

0000-0003-1435-3053]Richard W. Pogge
Department of Astronomy, The Ohio State University, 140 W 18th Ave., Columbus, OH, 43210
Center for Cosmology & AstroParticle Physics, The Ohio State University, 191 West Woodruff Avenue, Columbus, OH 43210

0000-0002-5659-4974]Simon Gazagnes
Department of Astronomy, The University of Texas at Austin, 2515 Speedway, Stop C1400, Austin, TX 78712, USA

0000-0002-0361-8223]Noah S. J. Rogers
Department of Physics and Astronomy, Northwestern University, 2145 Sheridan Road, Evanston, IL 60208, USA
Center for Interdisciplinary Exploration and Research in Astrophysics (CIERA), Northwestern University, 1800 Sherman Avenue, Evanston, IL 60201, USA

0000-0001-9714-2758]Dawn K. Erb
Center for Gravitation, Cosmology and Astrophysics, Department of Physics, University of Wisconsin Milwaukee, 3135 N Maryland Ave., Milwaukee, WI 53211, USA

0000-0002-2644-3518]Karla Z. Arellano-Córdova
Institute for Astronomy, University of Edinburgh, Royal Observatory, Edinburgh, EH9 3HJ, UK

0000-0003-2685-4488]Claus Leitherer
Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA

Jackie Appel
Associate News Editor, Popular Mechanics Magazine

0000-0002-2733-4559]John Moustakas
Department of Physics & Astronomy, Siena College, 515 Loudon Road, Loudonville, NY 12211

We investigate the stellar and nebular properties
of 9 H2 regions in the spiral galaxy M101 with far-ultraviolet
(FUV; $\sim900-2000$ Å) and optical ($\sim 3200-10000$ Å) spectra.
We detect significant C3] 1907,1909 nebular emission in 7 regions, but
O3] 1666 only in the lowest-metallicity region.
We produce new analytic functions of the carbon ICF as a function of
metallicity in order to perform a preliminary C/O abundance analysis.
The FUV spectra also contain numerous stellar emission and P-Cygni features that
we fit with luminosity-weighted combinations of single-burst  and  models.
We find that the best-fit  models closely match the observed very-high-ionization
P-Cygni features, requiring very-hot, young ($\lesssim3$ Myr),
metal-enriched massive stars.
The youngest stellar populations are strongly correlated with
broad He2 emission, nitrogen Wolf-Rayet (WR) FUV and optical spectral features, and
enhanced N/O gas abundances.
Thus, the short-lived WR phase may be driving excess emission in several N P-Cygni wind
features (955, 991, 1720) that bias the stellar continuum fits to higher metallicities
relative to the gas-phase metallicities.
Accurate characterization of these H2 regions requires additional inclusion of WR
stars in the stellar population synthesis models.
Our FUV spectra demonstrate that the $\sim900-1200$ Å FUV can provide a strong
test-bed for future WR atmosphere and evolution models.

§ INTRODUCTION

The history of a galaxy can be traced from its cumulative abundance pattern,
as abundances of heavy elements increase with time as successive generations of stars
return their newly synthesized elements to the interstellar medium (ISM).
Star-forming regions (H2 regions) are privileged sites to measure
the chemical abundances in galaxies, near and far, due to their ubiquitous
strong emission lines in the rest-frame optical regime.
While the majority of nebular studies come from optical observations,
the far-ultraviolet (FUV; defined here as 1200 – 2000 Å) is a rich wavelength
regime with many nebular diagnostic features that are complementary to
those in the optical.
Recently, the James Webb Space Telescope (JWST) has enabled rest-frame
optical and ultraviolet (UV) spectra of intermediate and high redshift galaxies,
ushering in a new era of chemical abundance studies across cosmic redshifts
Robust abundance measurements are key to deciphering
the complex physical and evolutionary processes of galaxy formation and growth.
In this regard, securing details of present-day abundances in star forming regions,
which better allow us to interpret the higher redshift spectra, is crucial for forming
the basis of our understanding of the chemical evolution of the universe.

Different elements trace different portions of chemical evolution owing to the
different times scales of their nucleosynthetic production.
On short time scales ($< 10$ Myr), massive stars ($>10\ M_\odot$) forge heavier elements
through the $\alpha$-process which fuses $\alpha$-particles into successfully more massive
elements to create oxygen, neon, magnesium, silicon, and beyond up to iron.
When these massive stars reach the iron-peak they explode as supernova and deposit
these $\alpha$-enriched elements into the ISM.
Carbon and nitrogen, crucial elements for the formation of dust and regulation
of ISM conditions, are not significantly produced by the $\alpha$-process.
Rather, the primary mechanism to produce nitrogen and carbon is thought to be neutron
capture within asymptotic giant branch (AGB) stars, occuring on longer timescales.
AGB stars are evolved lower mass (0.5–8 M$_\odot$) stars with inert cores of carbon
and oxygen.
These AGB stars are thought to be the primary source of nitrogen at low-metallicity
(12+log(O/H) $\lesssim$ 8.0; $Z\lesssim0.2\ Z_\odot$), with enrichment timescales
<cit.> that occur significantly after the most massive stars have exploded as

At higher metallicity ($Z\gtrsim0.2\ Z_\odot$), the carbon-nitrogen-oxygen (CNO)
cycle provides a “secondary" source of N, meaning that it is metallicity-dependent.
The CNO cycle uses trace amounts of carbon and oxygen as catalysts in a chain of reactions
that convert H to He.
Nitrogen is mostly produced in CN branch of the CNO cycle, which uses $^{12}$C as a
reaction catalyst and creates $^{14}$N as a byproduct.
The slowest “bottleneck" step of this branch is the $^{14}$N($p, \gamma$)O$^{15}$ step
because it has a relatively low cross-section, allowing
nitrogen to build up in the cores of high-metallicity massive stars.
Convection can then dredge up nitrogen to the surface, where stellar winds can advect it
into the ISM.
This process can enhance N-production during certain phases of massive star
evolution, such as the Wolf Rayet (WR) star phase, where stellar winds are
especially strong.
Measuring the relative chemical abundance (e.g. N/O or C/O) describes which of these processes a galaxy has undergoing and what the past star formation history of the galaxy looks like.

Recent observations of high-redshift galaxies have raised questions about the
standard view of chemical abundances.
For instance, deep rest-frame FUV observations of the galaxy GN-z11 find strong
N4] 1486, N3] 1750, and C3] 1907,1909 emission lines that
suggest that the galaxy has been enriched to super solar nitrogen abundances just a few
100 Myr after the Big Bang <cit.>.
This is far too rapid for enrichment by AGB stars.
Previous studies suggest that either WR, super massive stars, tidal disruption events,
or active galactic nuclei environments could be responsible for the high nitrogen abundances
However, the nitrogen and carbon lines observed within GN-z11 are rarely detected in local
star-forming galaxies, such that a complete understanding of the production and origin of
these features has not fully been revealed.
The origin of these features requires a multi-wavelength approach that incorporates
both the stellar and nebular physics.

UV spectra offer a unique glimpse into the processes occurring
in star-forming galactic environments.
In particular, the FUV regime
contains a rich diversity of features that can simultaneously
characterize the ionizing sources, feedback energetics, and
gas conditions within galaxies.
FUV stellar spectra of ionized regions (H2) trace the massive star population, as
they are dominated by P-Cygni stellar wind features, composed of a combination of broad,
blueshifted absorption and redshifted emission.
The so-called stellar wind P-Cygni profiles are produced by strong,
radiatively-driven stellar photosphere winds <cit.>.
Winds from O-stars are usually ionized to the N$^{+3}$, C$^{+4}$, and Si$^{+4}$
ionization states [Lamers et al., 1999].
However, none of these ionization states have strong resonant lines blueward of
Ly$\alpha$ (note that N4 1720 is a non-resonant line).
The lines that are available, N5 1238,1246, Si4 1393,1403,
and C4 1548,1550, trace adjacent ionization stages.
As such, these stellar wind signatures are highly sensitive to the temperature
(and, therefore, the age), density, and metallicity of the winds of the most
massive / luminous stars.
This sensitive dependency on stellar properties has important ramifications for
which lines trace which stars.
For instance, the Si4 1392,1402 doublet traces lower ionized gas
than the peak ionization, such that Si4 is typically only observed in
populations with cooler stars, such as evolved Giants or B-stars.
The N5 1238,1242 doublet on the other hand traces more highly ionized
gas and is only found in the hottest / youngest stellar populations.

In this paper we present HST/COS FUV spectral observations for 9
H2 regions in M101 in  <ref> and use these spectra to explore
the stellar features and nebular emission.
M101, also known as NGC 5457, is an invaluable target for investigating the stellar and
nebular properties across spiral galaxies.
Its utility stems from its proximity, allowing for detailed, spatially-resolved
investigations, and its face-on orientation, which presents an unobstructed view of its
spiral arms and associated nebulae, providing a comprehensive look at the distribution
and composition of stellar and nebular properties.
Our sample of H2 regions and spectral observations are described in
Sections <ref> and <ref>, respectively.
In Section <ref>, we perform stellar continuum fits to determine the
average properties of the massive star populations.
We discuss the FUV spectral characteristics of this sample in Section <ref>,
including the stellar continuum properties in  <ref> and
the presence of Wolf-Rayet features and He2 emission in  <ref>,
In  <ref> we examine the nebular C and O emission features and C ionization
correction factors (ICFs) from our grid of photoionization models in order to perform
an exploratory FUV C/O abundance analysis and compare to optical C/O measurements.
Our conclusions are given in Section <ref>.

§ THE M101 H2 REGION SAMPLE

For the present study, we targeted the star-forming regions identified by
Skillman et al., 2020 as very high surface brightness H2 regions in M101.
In particular, these H2 regions contain detections of $4-5$ intrinsically-faint
auroral collisionally-excited lines (CELs; see Table <ref>) and the even fainter
C2 4267 recombination line (RL) in their
high-quality optical spectra from the CHemical Abundances Of Spirals
<cit.> project study of M101 [Croxall et al., 2016].
In total, Skillman et al., 2020 presented 10 H2 regions.
For this work, we targeted the same sample of 10 H2 regions,
but one of the UV spectral observations failed, resulting in a sample of
9 H2 regions.

We note that Skillman et al., 2020 assumed the same properties for M 101 as in
Croxall et al., 2016, i.e., a distance of 7.4 Mpc [Ferrarese et al., 2000], resulting in
a spatial scale of 35.9 pc/arcsec; however, Beaton et al., 2019 have made a (hopefully)
definitive determination of the distance to M101 of 6.52 $\pm$ 0.12$_{stat}$ $\pm$ 0.15$_{sys}$ Mpc,
with a resulting scale of 31.6 pc/arcsec.
We have adopted the updated Beaton et al., 2019 distance here,
although updates to the distance to M101 will not affect the major conclusions
of this paper.
We have also assumed disk scalings of
R$_{25}$ $=$ 864 <cit.> and
R$_e = 198$ <cit.>,
an inclination angle of 18, and a major-axis position angle of 39 [Walter et al., 2008].

CHAOS uses the Multi-Object Double Spectrographs <cit.> on the
Large Binocular Telescope <cit.> to observe moderate-resolution
($R\sim2000$) optical spectra (3,200 Å $\lesssim \lambda \lesssim 10,000$ Å) of
a large numbers of H2 regions in spiral galaxies.
As part of the CHAOS program, very high quality spectra of 74 H2 regions in M101
[Croxall et al., 2016] were obtained that allowed direct determinations of absolute and relative
abundances across a broad range of parameter space.
As an extension of the CHAOS project, Skillman et al., 2020 examined the M101 spectra for
the optical RL and found detections of the C2 4267 line in 10 H2 regions.

Properties of M101 H2 Regions from the CHAOS Surveys

5c     5cAuroral Line Detections
Region      Name        R.A., Decl.     Location ID      $\frac{R}{R_{e}}$      12+log(O/H)
[O3]     [N2]     [S3]     [O2]     [S2]     He2     WR
1      H1013           14:03:31.3,+54:21:05.81     +164.6+009.9          0.89         8.60$\pm$0.03

2      H1052           14:03:34.1,+54:18:39.60     +189.2–136.3         1.30         8.58$\pm$0.02

4      NGC 5462        14:03:53.0,+54:22:06.80     +354.1+071.2          1.89         8.51$\pm$0.10

5      NGC 5462        14:03:53.8,+54:22:10.81     +360.9+075.3          1.99         8.51$\pm$0.02

6      NGC 5447        14:02:30.5,+54:16:09.89     –368.3–285.6        2.42         8.45$\pm$0.02

7      NGC 5455        14:03:01.1,+54:14:27.97     –099.6–388.0        2.19         8.39$\pm$0.01

8      NGC 5447        14:02:27.8,+54:16:25.14     –392.0–270.1        2.49         8.36$\pm$0.02

9      H1216           14:04:10.9,+54:25:19.20     +509.5+264.1          3.56         8.29$\pm$0.06

10     NGC 5471        14:04:29.0,+54:23:48.56     +667.9+174.1          3.81         8.16$\pm$0.02

Properties of the 9 H2 regions from the M101 sample that were observed with HST/COS.
Note that 10 regions were originally proposed, but the HST/COS observations of Region 3 failed.
Column 1 lists the region numbers that we adopt in this paper, whereas Column 2 gives the regions common name from either
the New General Catalogue (NGC) or the Hodge et al., 1990 catalogue.
Column 3 gives the right ascension (R.A.) and declination (Decl.) in J2000 and Column 4 gives the CHAOS ID from Croxall et al., 2016,
which is composed of the offset R.A. and Decl., in arcseconds, from the galaxy center.
Column 5 lists the galactocentric distance as a fraction of the effective radii, $R_e$, and Column 6 lists the oxygen abundances,
as updated by Berg et al., 2020.
Columns 7–11 highlight regions with significant auroral line detections of [O3] 4363, [N2] 5755,
[S3] 6312, [O2] 7230,7330, and [S2] 4069,4076.
We note in columns 12 and 13 which H2 regions have detections of narrow, nebular He2 4686 emission
and broad Wolf-Rayet features.

Image of M101 showing the H2 regions targeted by the CHAOS survey
using slit masks <cit.> and the 10 regions in which
recombination-line C2 4267 emission was detected <cit.>.
These 10 regions were followed-up with HST/COS observations as part of the
HST-GO-15126 program (PI: Berg).
The regions numbers correspond to those given in Table <ref>.

§ SPECTRAL OBSERVATIONS

§.§ LBT/MODS Optical Spectra

The optical observations presented here were previously reported in Croxall et al., 2016
and updated in Berg et al., 2020, where the interested reader can find details of
the observations and data reduction.
Here we briefly summarize.

Optical spectra of M101 were taken using MODS1 on the LBT during the spring semester of 2015.
We obtained simultaneous blue and red spectra using the G400L (400 lines mm$^{-1}$, R$\approx1850$)
and G670L (250 lines mm$^{-1}$, R$\approx2300$) gratings, respectively.
The primary objective of the CHAOS program is to detect the intrinsically weak
temperature-sensitive auroral lines in the wavelength range from 3200–10,000 Å
in order to determine direct abundances.
Although not a design goal of the program, the observations of M101 were sensitive enough
to detect RL emission from the C2 4267 line.
Recombination line observations are typically made at higher spectral resolution
than in the CHAOS program because O2 RLs are tightly clustered.
However, the 4267 line is well isolated, making C2 RL detections and abundances
of C$^{+2}$/H$^+$ possible from the MODS spectra.
Further discussion of the 4267 RL detections can be found in Skillman et al., 2020.

The original footprint of CHAOS observations of M101 is shown in Figure <ref>,
with the 10 regions from the current sample designated with green diamonds.
The present H2 region sample has a radial coverage of
$0.2\lesssim R/R_{25}\lesssim 0.8$ and spans a range of 0.44 dex in oxygen abundance.
Properties determined for the H2 regions in the present sample by the CHAOS survey
are reported in Croxall et al., 2016; a subset of important properties are listed in Table <ref>.

HST/COS G140L spectra of the 9 H2 regions in the CHAOS M101 sample,
normalized at 1800 and ordered by
increasing gas-phase oxygen abundance from top to bottom.
The location of potential UV spectral features are denoted for ISM absorption lines (green
long-dashed lines), stellar wind and photospheric features (gold dashed lines), and
nebular emission lines (black dashed lines).
The combination of photospheric lines at $\sim1430$ Å contains C3 1426,1428 and Fe5 1429,1430 in O stars <cit.>.
Note that the resonant transitions that produce stellar wind P-Cygni profiles
can also have ISM absorption components.
Due to the low-resolution of the G140L spectrum, MW features (locations indicated by gray
bands) are not resolved and may contaminate other features.
Note that there is also Geocoronal O1 emission near 1300 in the spectra
of Regions 4 and 5.

§.§ HST/COS UV Spectra

In order to measure the UV stellar and nebular properties of the H2 regions in M101
we obtained Cycle 25 HST/COS observations (program HST-GO-15126, PI: D. Berg).
Berg et al., 2016 demonstrated the utility of simultaneously observing the O3] 1661,1666
and C3] 1907,1909 nebular emission lines in the FUV afforded by the large wavelength coverage
(1282–2148 Å in Segment A) of the low-resolution G140L grating.
Adopting this observing strategy, we observed each of 9[All 10 H2
regions in the Skillman] et al., 2020 sample were proposed for in program HST-GO-15126,
however, the target acquisition failed on the final target (region 3, NGC 5461).
Since Space Telescope Science Institute deems a program as successful when 90% complete,
the observation was not repeated.
H2 regions in M101 for 1-3 orbits.

Due to the proximity of M101, the large H2 regions fill the COS aperture; however,
target image acquisition was efficiently achieved using a pointing offset from a nearby bright star.
We used deep B-band LBT Large Bincoluar Cameras (LBCs) imaging from
the OSU Monitor program <cit.>
to select bright, isolated stars near our target H2 regions.
The excellent astrometry of these images provided precise positions of the offset stars and
allowed us to accurately determine offsets.
We used the ACQ/IMAGE mode with the PSA aperture and MirrorB (due to the
brightness of the offset stars) with the COS/NUV configuration to acquire the
stars and then blind offset to the target H2 regions.

COS FUV spectral observations were taken in the TIME-TAG mode using the 2.5
PSA aperture and the G140L grating at a central wavelength of 1280Å.
We used the FP-POS = ALL setting, which takes four images offset from one another in
the dispersion direction, increasing the cumulative S/N and mitigating the effects
of fixed pattern noise.
The four positions allow a flat to be created and for the spectrum to fall on different
areas of the detector to minimize the effects of small scale fixed pattern noise.
Spectra from the 9 successful observing visits were processed with CALCOS version 3.3.4.
In order to gain signal-to-noise, we followed Berg et al., 2019 and re-binned the
spectra by a factor of six in the dispersion direction and smoothed with a 3-pixel
box-car kernel.
Nominally, the COS has a resolution of $R\sim2000$ for point sources, but this
resolution is degraded for nebulous regions that fill the aperture.
By measuring individual photosphere absorption lines in our spectra,
we found a typical FWHM$\sim300$ km s$^{-1}$, which corresponds to $R\sim1000$.

Comparison of Stellar and Gas Properties

6cStellar Properties                      2cGas Properties
3-8 10-11
Reg.            E(B-V)$_{\rm SB99}$     E(B-V)$_{\rm BPASS}$     Age$_{\rm SB99}$     Age$_{\rm BPASS}$     Z$_{\star,{\rm SB99}}$    Z$_{\star,{\rm BPASS}}$         O/H$_{neb.}$     N/O$_{neb.}$     $\Delta$Z$_{\rm SB99}$     $\Delta$Z$_{\rm BPASS}$
#       WR     (mag.)     (mag.)     (Myr)     (Myr)     (Z$_{\odot}$)     (Z$_{\odot}$)        (O/H$_\odot$)    (N/O$_\odot$)     (Z$_{\odot}$)     (Z$_{\odot}$)
1          0.218$\pm$0.003     0.178$\pm$0.004     1.94$\pm$0.03     3.65$\pm$0.56     1.00$\pm$0.03     0.88$\pm$0.01        0.81$\pm$0.01     0.81$\pm$0.01     0.19     0.06

2          0.197$\pm$0.004     0.172$\pm$0.003     1.60$\pm$0.04     2.76$\pm$0.19     1.00$\pm$0.03     0.88$\pm$0.01        0.78$\pm$0.01     0.74$\pm$0.01     0.22     0.10

4              0.239$\pm$0.035     0.258$\pm$0.020     4.12$\pm$2.33     4.15$\pm$2.38     0.49$\pm$0.13     0.51$\pm$0.07        0.66$\pm$0.02     0.44$\pm$0.04     -0.17    -0.15

5              0.245$\pm$0.014     0.251$\pm$0.015     2.42$\pm$0.64     5.57$\pm$4.03     0.89$\pm$0.11     0.47$\pm$0.04        0.66$\pm$0.01     0.43$\pm$0.01     0.23     -0.19

6          0.311$\pm$0.021     0.311$\pm$0.016     1.64$\pm$0.77     1.83$\pm$2.06     0.93$\pm$0.15     0.60$\pm$0.03        0.58$\pm$0.01     0.55$\pm$0.01     0.35    -0.03

7          0.195$\pm$0.005     0.182$\pm$0.005     2.42$\pm$0.11     4.30$\pm$1.00     0.78$\pm$0.05     0.52$\pm$0.02        0.50$\pm$0.01     0.55$\pm$0.01     0.28     0.02

8          0.174$\pm$0.005     0.159$\pm$0.005     2.26$\pm$0.09     2.40$\pm$0.96     0.99$\pm$0.04     0.55$\pm$0.02        0.47$\pm$0.01     0.59$\pm$0.01     0.52     0.08

9              0.169$\pm$0.011     0.192$\pm$0.011     3.31$\pm$2.16     7.14$\pm$3.40     0.62$\pm$0.15     0.28$\pm$0.07        0.40$\pm$0.01     0.33$\pm$0.03     0.22    -0.12

10             0.158$\pm$0.016     0.150$\pm$0.007     9.69$\pm$3.90     12.2$\pm$3.05     0.29$\pm$0.06     0.22$\pm$0.02        0.30$\pm$0.01     0.34$\pm$0.01    -0.01    -0.08
Stellar and nebular properties for the nine H2 regions studied here.
The stellar properties are derived from the luminosity-weighted stellar continuum fits
described in Section <ref>, where columns 3–8 correspond to the derived
reddening, age, and metallicity of the massive star population from both  and  fits.
The differences between the stellar model metallicities and the gas-phase metallicity are listed in
columns 9–10.
For comparison, regions with identified WR features are noted in Column 2, while the
nebular O/H and N/O abundances are given in columns 9–10.
All abundances are relative to the solar scale from Asplund et al., 2021, e.g., N/O$_{neb.}=0.813$ N/O$_\odot$
is equivalent to 81.3% the solar N/O value.

Figure <ref> shows the FUV spectra for the 9 H2 regions observed.
The observed wavelength coverage of G140L is $\sim900-1195$ for Segment B
and $\sim1282-2000$ for Segment A[Segment A coverage is quoted out to 2148 Å, however,
previous experience with G140L spectra suggest G140L
Segment A spectra only have utility out to $\sim2000$ Å.].
The 105 Å gap between Segments A and B mitigates the risk of strong Geocoronal Ly$\alpha$
emission on the detector, however, this also causes our spectra to miss the N5
1238,1242 stellar wind feature.
The spectra were corrected for the redshift of M101 ($z=0.0008$),
but not for the small offsets resulting from the galactocentric rotational velocity
of a given H2 region.

In order to allow a close comparison of feature trends, the spectra in Figure <ref> are
vertically offset from one another and ordered from top to bottom by increasing metallicity.
Nebular features are labeled by gray short-dashed lines, while stellar wind and photospheric
features are gold, and ISM absorption is labeled green.
Some features are complex combinations of multiple components, such as C4,
which can have stellar, nebular, ISM components, and, at the small redshift of M101,
contamination from Milky Way absorption.
Some features are clearly sensitive to metallicity, such as the S4, C4, and N4
stellar wind features, which strengthen with increasing metallicity, and the C3]
nebular emission, which generally weakens in strength at higher metallicity.
We also note the presence of He2 emission, which can be both stellar
and nebular in origin.
As we will see later, the strong, broad He2 features in regions 1, 2, 6, 7, and 8 are
likely stellar, as strong Wolf-Rayet (WR) features are noted in their optical spectra.
On the other hand, narrow He2 emission was found in the optical spectra of
regions 9 and 10, so their UV He2 features may be nebular (see Table 1).

§ STELLAR CONTINUUM FITTING

Luminosity-weighted best fit stellar models to the
HST/COS FUV spectra of the M101 H2 region sample.
Observed spectra are shown in turquoise, with the best-fit stellar
and  models overlaid in red and blue, respectively.
Rest-wavelengths of stellar wind and photospheric features are shown by
solid green and dashed gold lines, respectively.
Additionally, H and He lines are denoted by solid gray lines; the MW H lines are
indicated with dashed-dotted gray lines.
Fits were performed after masking any contamination from MW sky lines,
ISM absorption features, and nebular emission lines (gray-shaded bands).
The resulting age and metallicity that best characterizes the ionizing
stellar population is labeled in the upper right hand corner
of each panel for both the  (red) and
(blue) models.
We also highlight the excellent  WR model fit to the
N3 991 feature of Region 1 in the yellow inset box.
In general, there is a progression from strong stellar features at the
highest metallicities (top) to small features at the lowest
metallicities (bottom).

The interplay between gas and stars is one of the most fundamental,
yet unsettled, drivers of galactic baryon evolution.
The FUV ($\sim912-2000$ Å) is arguably the richest spectral regime in
diagnostic features characterizing the processes involved in the baryon feedback cycle.
This spectral range not only provides access to the only strong C nebular emission lines,
but also reveals stellar wind and photospheric features that are more sensitive to
the age and metallicity of the ionizing stellar population than any optical
diagnostic features <cit.>.
Consequently, the FUV continuum offers a valuable means to characterize the stellar clusters
within H2 regions, enabling a direct link between the properties of massive stars
and the nebular emission lines they power.
Further, in the high surface-brightness H2 regions targeted here, the
ionizing stellar population recently formed from the surrounding gas such that
the nebular and stellar chemical compositions are expected to be similar.
Here we take advantage of the HST/COS FUV spectra of M101 to model the
stellar continua of our H2 region sample and better understand the
conditions powering their observed emission-line spectra.

§.§ Fitting Method

We fit the stellar continuum following the procedure detailed in Chisholm et al., 2019.
This method uses stellar population synthesis modeling to characterize
the ages and metallicities of the ionizing stellar populations.
Fitting was performed as a linear combination of single-age, fully theoretical
stellar continuum burst models matched to the non-ionizing FUV continua of
our sample.
We use two sets of theoretical stellar models:
(1) single star models from  <cit.> and
(2) binary star models from  v2.2.1
With both set of models we employ a range of possible metallicities covering
0.05, 0.20, 0.40, 1.0, and 2.0 Z$_\odot$ and a range of burst ages (1–40 Myr),
use a high-mass cutoff of 100 $M_\odot$.
The best intrinsic spectrum is combined with a foreground dust screen model and
the Reddy et al., 2016 attenuation law to determine the reddening due to dust ($E(B-V)$).

We fit the entire observable range of the HST/COS G140L segment A spectra, covering $\sim1300-2000$ Å.
Next, we convolved the 0.4 Å resolution of the  models and the
1 Å resolution of the  models to the spectral resolution of the
individual H2 region spectra <cit.>.
In order to isolate the purely-stellar portions of the spectra, we mask any
features that are not produced by the stars, including Milky Way absorption
features and ISM emission and absorption lines from the H2 regions.
A list of features considered for masking can be found in [Leitherer et al., 2011].
The errors on the stellar fits were determined using a boot-strap Monte
Carlo procedure, where 300 iterations were performed
using a modified version of the observed spectrum with error randomly sampled
from a normal distribution.
Further details of the fitting method can be found in Chisholm et al., 2019.

The stellar continuum modeling simultaneously constrains the light-weighted age,
metallicity, and reddening ($E(B-V)$) of the ionizing massive star
population, the results of which are listed in Table <ref>.
We note that Regions 1 and 2 were best fit by  with super-solar light-weighted
metallicities that are significantly higher than their gas-phase abundances.
Therefore, we reran the  fits without the 2.0 Z$_\odot$ model,
effectively enforcing an upper metallicity limit of solar for these two H2
regions in order to more closely match the observed gas-phase oxygen abundance range.

The uncertainties reported from the Monte Carlo procedure for the reddening, age,
and metallicity of the stellar population are listed in Table <ref>.
In general, the uncertainties on the reddening and metallicity values are small
($\lesssim25$% for  and $\lesssim10$% for ).
However, a few of the regions have large age uncertainties, especially the regions
lacking WR features and regions older than 5 Myr.
These more uncertain ages are likely due to degeneracies arising from the lack of
strong spectral features able to distinguish older population ages.
Similarly, Sirressi et al., 2022 used integrated stellar population spectra with the
same continuum-fitting method presented here and compared to resolved star photometry
to show that ages could be measured with high accuracy ($<1$ Myr)
for populations with ages between 2 and 5 Myr.
This suggests that our youngest stellar ages may have underestimated uncertainties,
but they are robust within 1 Myr.

§.§  and  Comparison

The best stellar continuum fits using both  and
are shown in Figure <ref>, over-plotted on the observed HST
spectra, where the masked portions for each H2 region are denoted by
gray bands.
However, the  and  stellar continuum fits look very visually
similar for most of the H2 regions over the fit range (right column panels).
Comparing the results in Table <ref>, we find that the
fits have older light-weighted ages than the  models by 1.6 Myr
on average.
Countering this effect, on average, the  fits have 23% lower
metallicities than the  fits.
The fitted reddening values are similar between  and ,
with an average difference of only 0.006.

A visual comparison of the stellar features in the  and
models relative to the observed spectra can provide further insight.
In the fitted part of the spectra (right column panels), there are three stellar
wind features – Si4 1393,1403, C4 1548,1550, and
N4 1720 – whose wind profiles are all fit reasonably well by both
and .
However, excess Si4 and N4 emission are seen in several
of the WR regions.
Additionally, neither  nor  have WR prescriptions that
can produce sufficiently broad He2 emission in these regions.

While not fit with the stellar models, we also examined the bluest portion of
the COS spectra.
The $\sim900-1200$ Å portion of the FUV is notoriously difficult to fit due to
the densely packed suite of stellar and ISM lines in the observed galaxy and
low-ionization resonance lines formed in the intervening gas, including numerous
Lyman series H lines and Lyman and Werner bands  of H$_2$ <cit.>.
This leaves little to no uncontaminated continuum to fit <cit.>.
Therefore, it is interesting that both the
and  models produce good overall fits to the observed spectra,
despite being unconstrained by the data in this challenging wavelength regime.

On the other hand, the fit of the  models to the most prominent features
in the bluest spectral regime is rather impressive.
In particular, the  models trace and confirm the O6 1032,1038
P-Cygni profile in Regions 1, 2, 5, 6, 7, and 8.
The O6 feature is further evidence of very young stars in these regions ($<3$ Myr; see next
section), favoring the younger ages of the  fits.
Note, however, that the  models do not incorporate X-ray winds,
and so they will not be able to fit the O6 (I.P. $\gtrsim 114$ eV) feature.
It would also be useful to fit the age-sensitive N5 1238,1242 feature in
order to confirm the young ages of the massive star population, but N5
unfortunately falls in the G140L gap (implemented to avoid Geocoronal Ly$\alpha$) at
the low redshift of M101.

Comparison of the O4, C4, and N4 stellar features in
Regions 1 and 2 for  models with different light-weighted metallicities.
Regions masked for possible contamination from MW sky lines or
ISM absorption features are indicated by gray-shaded bands.
There is very little difference between the two sets of models, and the models
fail to reproduce the peak of the emission in the N4 P-cygni profile,
demonstrating the need for both improved models and higher resolution and S/N spectra.

§.§ Single- and Mutli-Population Comparison

Typically, simple stellar populations (SSPs) assuming a single-burst (single-population)
are used to characterize the stellar continuum of an H2 region, while multiple
single-burst models (multi-population) are combined to characterize the integrated spectra
of galaxies.
As demonstrated for Region 10 (NGC 5471) in Figure <ref>,
many of the H2 regions in M101 are giant complexes containing multiple knots of
star-formation and earning them their own NGC classifications.
For example, García-Benito et al., 2011 performed a photometric study of the star formation
history of NGC 5471 showing that star formation has been more or less continuous for
at least 100 Myr.
Further, they found that both the integrated photometry and resolved-star color-magnitude
diagram of NGC 5471 is consistent with star formation of two different ages.

More recently, Sirressi et al., 2022 used Legacy ExtraGalactic UV Survey (LEGUS) HST multiband
photometry and CLusters in the UV as EngineS (CLUES) FUV HST/COS spectroscopy of 20
young stellar clusters hosted in 11 nearby star-forming galaxies to examine their
stellar populations.
Using single-population fits, two-population fits, and the same multi-population fit method
used here, Sirressi et al., 2022 found that most of the star-forming regions examined
(16/20) are populated by two or more stellar populations.
In fact, a second population was required for these regions to accurately reproduce the
observed level of spectral continuum and the P-Cygni line profiles simultaneously.
Given the complicated nature of the giant H2 regions in M101 and their complex
star formation histories, we adopted the multi-population method to characterize their
stellar continua in our analysis above.

To test the validity of our multi-population assumption, we also fit the two regions showing
the strongest stellar features (Regions 1 and 2) with single-population fits.
The resulting single-population  and  fits (dashed-lines) are shown in Figure <ref>
in comparison to their corresponding multi-population fits (solid lines).
In general, the multi-population fits better fit the stellar wind profiles.
For Region 1, the single- and multi-population  fits show little difference, but the
multi-population  fit more closely matches the Si4 1393,1403 and C4
1548,1500 outflow profiles than the single-component  fit.
For Region 2, both the  and  fits show differences between the single- and
multi-population fits.
In particular, the multi-population fits better reproduce the Si4 1393,1403,
C4 1548,1500, and the N4 1720 profiles.

The resulting stellar continuum properties for the single- and multi-population  and
fits are compared in Table <ref>.
In general, the reduced $\chi^2$ values are similar between the single- and multi-component
fits, but are smaller for the multi-component fits.
Additionally, the multi-component fit properties have much smaller uncertainties relative
to the single-component fits.
However, an important aspect of the multi-population  fits is that they allow for
slightly older ages.
We, therefore, conclude that there are not significant differences between the single-population
and multi-population fits, but that the multi-population fits provide a small improvement
in the complex H2 regions of M101.

Comparison of Single- vs. Multi-Population Stellar Continuum Properties

3cReg 1        3cReg 2
2-4 6-8
Single     Multi.                     Single     Multi.
Property       Pop.       Pop.       $\Delta$        Pop.       Pop.       $\Delta$
Z$_{\star,{\rm SB99}}$ (Z$_\odot$)      1.000±0.238     1.000±0.028     +0.000        1.000±0.238     1.000±0.030     +0.000

Z$_{\star,{\rm BPASS}}$ (Z$_\odot$)     1.000±0.238     0.876±0.009     -0.124        1.000±0.238     0.881±0.006     -0.119

Age$_{\rm SB99}$ (Myr)                  2.000±2.000     1.940±0.033     -0.060        1.000±2.000     1.600±0.040     +0.600

Age$_{\rm BPASS}$ (Myr)                 1.000±2.000     3.650±0.558     +2.650        1.000±4.000     2.758±0.192     +1.758

E(B-V)$_{\rm SB99}$ (mag.)              0.247±0.033     0.218±0.003     -0.029        0.208±0.028     0.197±0.004     -0.011

E(B-V)$_{\rm BPASS}$ (mag.)             0.223±0.021     0.178±0.004     -0.045        0.180±0.010     0.172±0.004     -0.008

Reduced $\chi^2_{\rm SB99}$             14.63             14.22             -0.41         8.92              8.89              -0.03

Reduced $\chi^2_{\rm BPASS}$            9.72              8.27              -1.45         5.86              5.73              -0.13

Comparison of the properties derived for the single- and multi-population  and
stellar continuum fits for Regions 1 and 2.
The single-population fits tend to drive metallicities higher and ages lower relative
to the multi-population fits.
The multi-population fits also tend to have smaller reduced $\chi^2$ values.

This image of Region 10 (NGC 5471) using the F656N filter on HST/WFPC2
(PID: HST-GO-6829; PI: You-Hua Chu)
shows multiple knots of H$\alpha$ emission composing the complex, giant H2 region.
This suggests that multiple star associations are contributing to the gas ionization.
For reference, the white box spans 10$\times$10.

Comparison of the single-population best-fit stellar models (dashed lines)
to the luminosity-weighted multi-population best-fit stellar models for Region 1
(top panel) and Region 2 (bottom panel).
Observed spectra are shown in turquoise, with the best-fit stellar continuum models of
(red) and  (blue) overlaid.
Differences between the single- and multi-population fits are emphasized by red and blue
shading for the  and  models, respectively.
(blue) models.
As in Figure <ref>, stellar wind (solid green lines), photospheric features
(dashed gold lines), and masked regions (gray-shaded bands) are shown.
Overall, the multi-population fits better reproduce the observed stellar wind features,
and are adopted for the analysis in this work.

§ CHARACTERISTICS OF THE FUV SPECTRA OF M101

§.§ Stellar Continuum Properties

Figure <ref> shows stellar-wind P-Cygni profiles of O6 1032,1038,
Si4 1393,1402, C4 1548,1550, and N4 1720.
Since more massive and luminous stars have shorter main sequence lifetimes, the
P-Cygni profile strengths constrain the age of the ionizing stellar population.
Higher ionization state profiles, such as O6 and N5, are
strongest in stars with lifetimes of 2–3 Myr, while lower ionization state
profiles, such as C4 and Si4, are more prominent in stars with
lifetimes near 5 Myr.
This trend can be seen in Figure <ref>, were the C4 profile is generally
stronger for the lowest stellar population ages and weak for the highest ages.
This is further evidenced by Regions 1, 2, 6, and 7 having the strongest O6 P-Cygni
profiles and the youngest  ages.
Note, however, the O6 wind profile is contaminated by strong Ly$\beta$
absorption or emission, and so requires careful evaluation.

P-Cygni stellar-wind profiles are also sensitive to the metallicity of the stellar
photospheres, where continuum photons are absorbed and drive gas from the stellar surface.
However, this metallicity sensitivity declines above solar.
In Figure <ref>, we show the O4, C4, and N4 stellar features
for H2 regions 1 and 2 in our sample, both of which were initially best fit with a
super-solar luminosity-weighted stellar continuum metallicity.
In comparison, we show their respective adopted best-fit $Z=Z_\odot$ stellar continuum
models (red) and a super-solar $Z=1.45\ Z_\odot$ model (cyan).
In both regions, the O4 1341 photospheric feature is well fit by both models
blueward of the deep C2 1335 ISM absorption feature.
The C4 1548,1550 wind profiles are also well fit, but the models slightly
underfit the depth and height of the P-Cygni profiles, regardless of metallicity.
For the N4 1719 wind feature, the absorption component of the P-Cygni profile
is well fit blueward of the Ni2 ISM absorption feature, but both models fail to
reach the strength of the emission component.
However, an additional N4 source may be contributing to this profile,
as both Regions 1 and 2 appear to have N4] 1487 emission in Figure <ref>.

In Figure <ref>, we compare the derived stellar population metallicity versus
the nebular metallicity, as derived from the collisionally-excited lines direct method.
As expected, the stellar and nebular metallicities generally correlate.
Interestingly, the  stellar metallicities are generally consistent with the
nebular metallicities, while the  stellar metallicities are offset to higher
values, especially for the youngest-aged stellar populations.
We fit a linear relationship to both the  and  trends using the
Bayesian linear regression code
which is an implementation of the linear mixture model algorithm developed by
Kelly, 2007 to fit data with uncertainties on two variables, including explicit
treatment of intrinsic scatter.
The  fit, shown in orange, indicates that, on average, the stellar
metallicities are offset to higher metallicities by 0.25 dex.
Figure <ref> also shows a cluster of  outliers at the highest stellar
metallicities considered.

Taken as a whole, to match the strong stellar wind profiles of the youngest stellar
populations that contain observed WR profiles (red diamonds in Figure <ref>),
models require very young populations ($<3$ Myr) that are more metal-enriched
than implied by their gas-phase metallicities.
This may be because the lower-metallicity stellar evolution tracks do not produce significant
amounts of evolved stars [Leitherer et al., 2011].
Therefore, the code requires younger and more metal-enriched stellar populations to match
the strong profiles in the populations with observed WR stars.
Meanwhile,  fits do produce sufficiently strong wind profiles that mimic the
observed profiles at metallicities that are consistent with the observed gas-phase metallicities.
This may be because the binary synthesis in  allows for additional
physical pathways to create the evolved (WR) stars observed in the spectrum [Eldridge et al., 2017].
While  does fit some of these lines better, many of the WR features are still
poorly fit by population synthesis models.
Thus, these stellar populations provide a strong test-bed to constrain future WR atmosphere and
evolution models, and their incorporation to stellar population models.

Comparison of the derived stellar population metallicity versus
the nebular metallicity, as derived from the CEL direct method.
As expected, the stellar and nebular metallicities generally increase in lock step,
however, stellar metallicities are offset to higher values than nebular metallicities
for the youngest-aged ionizing stellar populations.
These outliers also tend to show WR features in their optical spectra (red diamonds),
suggesting the WR phase is strongest for young ($t < 3$ Myr),
metal-rich ($Z>0.9Z_\odot$) massive stars.

§.§ He2 Emission

Emission from He$^{+2}$ recombination produces the He2 emission features commonly
observed at 1640 in the rest-frame FUV and 4686 in the rest-frame optical.
He2 emission is expected from a number of sources, including shocks, X-ray binaries,
extremely metal-poor massive stars, and the winds of massive stars, such as WR stars.
WR stars are an interesting source because they are extremely hot stars that
have main-sequence lifetimes $\lesssim5$ Myr <cit.> being
observed during a short-lived supergiant phase.
During this phase, WR stars produce broad He2 emission lines in both the optical and
FUV due to their strong winds that increase at higher metallicity <cit.>.

Chisholm et al., 2019 used light-weighted stellar populations to show that the He2
profile is also very sensitive to age.
For the youngest populations in their study ($\lesssim 5$ Myr), He2 1640
has a strong, broad profile, while the oldest populations ($> 20$ Myr) show a weak
He2 absorption feature.
They also found that their multiple-age fits to the FUV He2 1640 region were
poor, especially in the presence of strong WR emission, with
performing worse than .
Similar to the results of Chisholm et al., 2019, Figure <ref> shows that our best-fit
stellar models fail to produce the requisite number of He2 photons from WR stars to
match the broad He2 emission observed in our HST/COS spectra of
H2 regions in M101, especially for the youngest stellar populations
Further, Martins et al., 2023 used  models with an initial mass function upper
mass cutoff of 300 $M_\odot$ to show that not even the winds of very-massive stars could
reproduce the broad strong He2 observed here.

§.§ Wolf-Rayet & Very Massive Star Features

§.§.§ Optical Characteristics

Given the low spectral resolution of the FUV spectra ($\gtrsim 300$ km s$^{-1}$),
the broad He2 emission at 1640 is difficult to interpret.
Therefore, it is useful to look to the optical spectra ($\Delta v \sim 150$ km s$^{-1}$)
to determine whether the He2 emission appears narrow and nebular in origin or
is broad due to WR or very-massive star ($>100\ M_\odot$; VMS) winds.
In Figure <ref>, we plot the He2 portion of the FUV spectra in comparison to
the two optical regions that are commonly used to spectroscopically identify the presence
of WR stars.
The first optical feature, known as the blue bump (middle column), is a complex of features
near 4650, including N5 4603, N3 4634,4641, C3 4647,4666,
C4 4658, and He2 4686.
The second feature is a blend of C4 transitions at 5801,5812, known as the red
bump (right column).
The shapes of these optical bumps help divide WR stars into three main classes
(1) WN stars: have H absorption lines and strong He and N emissions lines,
especially in the blue bump, but can also have a few C lines;
(2) WC stars: have strong He, O, and C emission lines, especially the C4 lines
characteristic of the red bump, but lack N lines; and
(3) WO stars: have strong lines from He and high-ionization lines from C and O.

Using the FUV and optical spectra presented here, we can characterize the
WR nature of the stellar populations in our sample.
We see that Regions 1 and 2 have clear N2 4621 and N3 4634,4641 emission in
their blue bumps, along with weak detections of the N4 4057 feature.
Weak N features may also be present in Regions 6, 7, and 8.
We also look for features indicative of WC stars.
Just redward of the N emission in the blue bump (and blueward of [Fe3]),
a second less prominent emission feature is seen in Regions 1 and 2, but it is difficult to
determine whether this is due to C3 or O2 emission, or both.
Further, the red C4 bump doesn't appear strongly in any of the H2
regions in our sample, suggesting little to no contributions from WC stars in our spectra.
Therefore, in the five regions with strong or weak optical WR features,
we only significantly detect the blue WR bump, suggesting that WN stars are likely present.
Using the Smith et al., 1996 N5 4057/N3 4603-4641 classification scheme,
we find these stars are most consistent with late WN6-WN8 stars in Regions 1 and 2.

On the other hand, Martins et al., 2023 recently examined optical spectra of six of the
H2 regions presented here for signatures of
VMS or WR features, but did find contributions from WC stars.
However, Martins et al., 2023 used optical spectra from the SDSS and from GTC/OSIRIS
from Esteban et al., 2020, which have slightly different aperture pointings, sizes, and
position angles than our CHAOS optical spectra.
Because these H2 regions have experienced intense star formation, they may
contain a mixture of massive star types that are spatially distributed and so not
equally captured in the different apertures used.
In the present study, we took special care to design our HST/COS UV
observations to have the same precise coordinates as the CHAOS optical spectra by
employing target offsets from nearby point-sources.
Therefore, the CHAOS optical spectra represent the closest spatial map to the
FUV HST/COS spectra.
However, there may still be differences in the stellar light captured between
the FUV and optical spectra owing to the different aperture shapes and the
vignetting of the COS aperture.
These differences in observational setup may also be the source of discrepancies
in identifying WR and VMS features.
For example, <cit.> observe both
the blue and red WR features in Region 7 (NGC 5455) and suggest that both WN and WC
stars are present, while our spectra show no definitive signatures of WC stars.

§.§.§ FUV Characteristics

In agreement with past studies, the prominent WR emission in the M101 optical spectra
corresponds to FUV spectra of young, metal-rich stellar populations (see Table <ref>).
Interestingly, the M101 observations presented here suggest a threshold for strong WR
emission of stellar $Z \gtrsim 0.78\ Z_\odot$ ($Z\gtrsim 0.54\ Z_\odot$) and
$t \lesssim 3$ Myrs ($t \lesssim 5$ Myrs) for the  ()
models; slightly younger and less metal-rich than previously thought
This is shown in Figure <ref> where we compare the stellar metallicities and
ages of the full H2 region sample to that of the WR subsample.
H2 regions with WR features only occupy the upper left young-age,
high-metallicity regime of the plot.

Such young ages for WR stars, as indicated by the  models, are somewhat
surprising given that the main sequence H-core burning phase of massive stars that
precedes the WR phase typically lasts longer than 3 Myr
However, the situation is further complicated by the fact that massive star and
WR evolutionary tracks are metallicity-dependent.
For example, Leitherer & Heckman, 1995 used evolutionary synthesis models for populations of
massive stars to show that WR stars can appear after just 2.5 Myr in twice-solar populations
and after just 2.0 Myr in solar populations.
Eldridge et al., 2017 also examined the presence of WR stars by subclass and found that
WC stars appear after 2.5 Myr in solar populations, in agreement with Leitherer & Heckman, 1995,
but that WN stars can appear slightly earlier after 2.0 Myr.
If WR stars can appear after only 2 Myr, then, within the uncertainties, the stellar
population ages for our sample allow for the presence of WRs.
Only the  age for Region 2 would be exceptional, but is also within agreement
assuming an accuracy of $<1$ Myr [Sirressi et al., 2022].

Alternatively, we could be observing a pre-WR phase of VMSs.
The characteristic emission features used to identify WR stars arise in their dense,
high-velocity stellar winds.
Because VMSs also produce exceptionally strong stellar winds and have high mass loss rates,
their spectra can mimic those of the later evolutionary stage of WRs
In general, the FUV spectra of VMSs and WRs are difficult to tell apart.
Both VMSs and WRs have been observed to have  many of the features present in our
HST/COS spectra, including N4] 1486, He2 1640, and
N4 1720 <cit.>.
On one hand, Martins et al., 2023 argue that the UV alone is insufficient and the optical WR
feature morphologies are needed to distinguish between VMSs and WRs.
As such, the presence of the blue bump in some of our spectra favors the presence of WR stars.
On the other hand, VMSs are only present at very young stellar population ages
($<2.5$ Myr), suggesting VMSs could be the source of the excess wind profiles in our
spectra and the corresponding young ages of the  fits.

The ionizing stellar populations in M101 could contain a complicated mixture of massive
stars with a range of ages ($<10$ Myr) such that both VMSs and WRs are present.
However, the presence of VMSs alone is deemed unlikely.
First, the presence of Si4 P-Cygni profiles in Regions 1 and 2 in Figures <ref>
and <ref> require older ages than what is physical for VMSs
Additionally, the presence of O5 1371 is considered a key spectral diagnostic
of VMSs, but there are no strong detections in the M101 spectra.
The observed spectrum of Region 1 in Figure <ref> shows weak O5 absorption,
indicating that VMSs may also be present in this region.
However, the other regions show either no absorption at O5 1371 or
blue-shifted absorption that is more consistent with Fe5 1365.
For the other WR spectra (Regions 2, 6, 7, and 8) in Figure <ref>, the ambiguous
absorption feature near 1370 is well fit by the Fe5 feature in the stellar
continuum model.
Therefore, WRs seem to be the dominant wind source in our spectra, although we cannot
Further, such large, complex H2 regions as are seen in M101 might be
likely places to find multiple, unresolved stellar populations hosting both WRs and VMS.

Proving the presence of strong WR stars at moderate metallicities is also challenging
because WR features have not been observed in the spectra of individual metal-poor
massive stars.
This may simply be due to the fact that metal-driven winds in the photospheres of
massive stars decrease with lower metallicities by definition.
As a result, WR features have also proven difficult to produce in theoretical models
at low  metallicities, while WR features appear at relatively high metallicities and
progressively increase in strength (e.g., broad He2 emission) in the stellar
population synthesis models.
Therefore, the presence of WR features in our observed spectra may be driving our
best continuum fits to higher luminosity-weighted metallicities.
In turn, requiring high metallicities to fit the WR features may drive the
luminosity-weighted ages younger.
This could explain the large offset in  stellar metallicity in
Figure <ref> for the subset of H2 regions with WR features in their spectra.

The WN-type stars in our M101 H2 region sample are further confirmed by their
relatively large N/O abundances in Table <ref> (larger for WR regions than
non-WR regions) and their FUV N spectral features.
As mentioned in Section <ref>, WR stars may also be responsible for the
N4] 1483,1487 emission and excess emission in the N4 1719 P-Cygni
profile are seen in Figures <ref> (also see Figure <ref>).
Such N4] detections are very intriguing.
First, N4] 1483,1487 emission is rare and has only been reported
for a handful of galaxies and active galactic nuclei (AGN) across all redshifts.
Second, there has been heightened interest in N4] due to its strong detection
in the JWST/NIRSpec spectrum of the exceptionally bright $z=10.6$ galaxy
GN-z11 [Bunker et al., 2023].
As a result, WR stars have been suggested as a source of increased N emission for
GN-z11 and other galaxies <cit.>.

The WR-driven N features seen in our M101 spectra are further supported by the presence
of P-Cygni profiles of the resonance or excited transitions of metal lines in bluest
portions of the FUV spectra, including S6 933,945, N4 955,
N3 991, O6 1032,1038, P5 1118,1128, and C3 1175.
This combination of features and range of moderate- to high-ionization potential species
is consistent with the later WN6-WN7 FUSE spectra described in Willis et al., 2004.
Additionally, if WC stars were present, we'd expect to see strong emission or P-Cygni
profiles of C3 977,1140 and/or C4 1107,1135,1169, yet we
see no evidence of significant WC stars in the blue segment of our FUV spectra.

§.§ Impact of WR Stars on Inferred Stellar Properties

While our integrated FUV spectra are difficult to fit with current stellar models
owing to the complex blend of massive star and WR contributions, it is clear
that the WR features can have a significant impact on stellar continuum and
subsequently derived properties.
The impact of WR stars on integrated star-forming spectra can be seen in the
residuals to our stellar population model fits.
The spectral regions where our stellar population synthesis fits perform the
poorest correspond to underfit WN P-Cygni and emission profiles (e.g., N4 1720),
suggesting that the
and  models fail to reproduce the full WR
populations in these observed regions.
Fortunately, the Potsdam Wolf-Rayet <cit.>
model atmospheres for WR stars can provide significantly better fits to the
FUV WN features.
In the yellow window of Figure <ref> we expand the observed spectrum of Region 1
around the N3 991 feature and overplot a stellar model from the
MW WNL-H20 grid that corresponds to $T_\star = 50.1$ kK and log $(R_t/R_\odot) = 0.3$.
After a relative scaling between spectra, it is clear that the N3 991
emission is easily reproduced by a WN star.
Therefore, the observed spectra require a combination of ordinary massive stars
and WN stars to full fit it.

On the other hand, the WN features discussed in this work can provide
powerful tools to improve our interpretation of FUV continuum spectra.
Using our M101 FUV spectra, we have demonstrated that the $\sim900-1200$ Å
stellar continuum features provide strong diagnostics of the presence of WR stars
that can be used to interpret the integrated spectra of star-forming regions or galaxies.
Further, these very blue spectra provide support for the presence of WN stars
as the source of excess N4 1483,1486 and N4 1720 emission.
As discussed above, the presence of very high ionization stellar wind profiles
suggests the presence of very-hot / young massive stars, whose integrated spectra are
better characterized by the  and  models.
However, the  and  models either do not have sufficient numbers of WR stars
or do not include appropriate WR evolutionary tracks to reproduce the excess
N emission observed in our spectra.
As a result, in trying to fit this excess emission,
the  model fits are biased to higher-metallicities and younger ages.
Therefore, the H2 regions of M101 provide excellent benchmarks for
constraining the revised inclusion of WR and VMS stars in stellar population models.
As such, we provide the emission-line properties for important UV and optical WR
features measured for Regions 1 and 2 in Table <ref>.

rcc c cc

Stellar Wind Emission-Line Properties

2cRegion 1        2cRegion 2
2-3 5-6
EW        FWHM                 EW        FWHM
Line       (Å)     (km s$^{-1}$)        (Å)     (km s$^{-1}$)
N4] 1486         $-1.37$     1434        $-0.72$     862

He2 1640         $-2.77$     1436        $-3.08$     1344

N4 1720          $-1.08$     603         $-1.54$     636

? 1726                   $-0.77$     603         $-0.81$     636

N3 4634          $-0.99$     950         $-1.78$     936

N3 4640          $-2.00$     1085        $-2.36$     748

C3 4650          $-0.33$     345         $-0.57$     294

He2 4686         $-3.20$     975         $-8.99$     1586

C4 5801          $-0.55$     775         $-0.27$     362

C4 5808          $-0.55$     775         $-0.38$     362

2cHe2 1640
2cHe2 4686
2-3 5-6
EW          FWHM        EW          FWHM

(Å)       (km s$^{-1}$)        (Å)     (km s$^{-1}$)
Region 1     $-2.77$     1436        $-3.20$     975

Region 2     $-3.08$     1344        $-8.99$     1586

Region 4     –          –          –          –

Region 5     $-1.81$     742         –          –

Region 6     $-2.05$     1113        $-3.73$     1202

Region 7     $-2.62$     1981        $-3.52$     1574

Region 8     $-3.34$     1666        $-3.36$     1201

Region 9     $-1.02$     859         $-0.68$     289

Region 10    $-1.95$     1420        $-1.31$     296

Emission-line properties of stellar wind features observed in the UV and optical
spectra of the M101 sample.
Line fits for the regions with the strongest wind features (Regions 1 and 2).
Column 1 lists the ion and wavelength of each emission line, with the exception of
the emission feature near 1726 that has an uncertain association.
Columns 2 and 3 list the equivalent width (EW) and Gaussian full width half maximum
(FWHM) for Region 1, while Columns 4 and 5 list EW and FWHM for Region 2.
UV and optical He2 emission properties for the full sample.
Note that the He2 features are not well fit in the low-resolution UV spectra
for regions with weak emission (EW$_{\lambda1640} < 2$Å).

Comparison of FUV He2 emission features (left panel)
with the optical WR blue bump (middle panel) and red bump (right panel) features.
We detect strong blue WR features in several of the regions, indicating the presence of
nitrogen-type WR (WN) stars.
Regions identified by Croxall et al., 2016 to have narrow, nebular He2
emission or WR emission in their optical spectra are labeled accordingly in
the upper right-hand corner of the left-panel plots.
We do not detect any red bumps in the right panel, and thus lack explicit evidence of
carbon-type WR (WC) stars, but there may be C3 contributions to the blue bump in the middle panel.

Top: Modeled luminosity-weighted metallicity versus age
for the ionizing stellar populations in our H2 region sample.
We show points derived using both  (circles) and
(semi-transparent squares) models.
Regions with WR features in their optical spectra are outlined by a red
diamond; these regions also have N4 emission in their FUV spectra
(turquoise points).
H2 regions with WR features only occupy the upper left young-age,
high-metallicity regime of the plot.
Bottom: N/O abundance measured from the optical spectra versus modeled stellar age
shows that the same regions that have enhanced N/O gas abundances have enhanced N4]
emission and strong WR features in their spectra.
This supports the idea that the WR phase is responsible for the observed local N
enhancement in these H2 regions.

Continuum-subtracted HST/COS spectra of the seven
H2 regions with C3] 1907,1909 emission.
Monte-Carlo minimization fits to the Gaussian profiles of O3] and
C3] are shown as black lines.
Many of the O3] profiles lack sufficient signal due to the neighboring
Al2 1671 ISM (orange dashed line) and MW (green dashed line)
absorption profiles.

§.§ O3] and C3] Emission

The FUV O3] 1661,1666 and C3] 1907,1909 nebular emission lines
offer the opportunity to investigate the relative C/O abundance in galaxies.
In general, carbon abundance determinations in extragalactic H2 regions are rare
in the literature due to the lack of bright CELs of C
in the optical and the extreme faintness of the optical RLs.
Previous C/O studies have been most successful in measuring the UV C and O CELs in
metal-poor dwarf galaxies <cit.>,
while optical RL studies dominate in metal-rich regions
(12+log(O/H) $>$ 8.0) of very nearby galaxies
Given this apparent dichotomy in the previous C/O measurements, comparisons between the
CEL and RL methods for C/O have remained sparse.

measurements of abundances in metal-rich H2 regions (12+log(O/H) $>$ 8.0) typical
of spiral galaxies are difficult in general.
Since the cooling in H2 regions is dominated by IR fine-structure lines of N$^{++}$
and O$^{++}$ for high metallicities, the greater the metallicity, the greater the cooling.
The lower electron temperatures mean that the higher excitation CELs become progressively
weaker, and the electron temperature becomes difficult to measure.
This is a problem especially for C, as the main ionization state, C$^{++}$, has its
strongest emission in the UV, and this becomes vanishingly faint in metal-rich nebulae.

Due to these challenges, there is a paucity of metal-rich targets with
significant UV C3] and O3] detections.
For example, Garnett et al., 1999 observed C3] in 4 metal-rich H2 regions
in the spiral galaxies NGC 2403 and M101 using FOS/HST, but were unable to obtain any
significant detections of the O3] doublet.
Because COS is significantly more sensitive than FOS, HST is now better able to detect
the C3] and O3] UV emission lines in metal-rich H2 regions,
such as the two regions with 12+log(O/H) $>$ 8.2 reported in Senchyna et al., 2017. =-2

§.§.§ Flux Measurements

The O3] 1661,1666 and C3] 1907,1909 emission line features of our
UV HST/COS spectra are plotted in Figure <ref>.
To determine the emission line strengths for these UV features, we first continuum-subtracted
the spectra using the best-fit stellar continua described in Section <ref>.
The UV emission lines were then fit using the  [Virtanen et al., 2020]
to perform a least-squares minimization on the built-in Gaussian function in
the   package [Newville et al., 2018].
Groups of nearby lines were fit simultaneously, constrained by a single Gaussian
FWHM and a single line center offset from the vacuum wavelengths (i.e., redshift).
Figure <ref> shows the Gaussian fits to the significant O3] and C3]
emission lines in the COS spectra of M101.

The errors of the flux measurements were approximated using

\begin{equation}
\sigma_{\lambda} \approx \sqrt{(2 \times \sqrt{N} \times \mbox{RMS})^2 + (0.01 \times F_{\lambda})^2},
\label{eq1}
\end{equation}

where N is the number of pixels spanning the Gaussian profile fit to the emission lines.
The root mean squared (RMS) noise in the continuum was taken to be the average of the RMS on
each side of an emission line when free of absorption features.
The two terms in Equation <ref> approximate the errors from continuum subtraction and flux calibration,
where the flux calibration uncertainty is assumed to be 1% of the flux based on
the standard star calibrations in CALSPEC [Bohlin, 2010].
For weak lines, such as the UV CELs, the RMS term determines the approximate uncertainty.
Finally, UV fluxes and errors were corrected for extinction using the reddening law of
Reddy et al., 2016 and nebular color excesses ($E(B-V)$) derived from the optical Balmer lines
from Croxall et al., 2016.
The reddening corrected line intensities measured for our C/O targets are reported in
the top of Table <ref>.

Seven of our 10 targets have significant C3] detections,
which we consider a high success rate given the relatively large O/H abundances of our sample.
However, only one region has a significant O3] detection.
For nearby galaxies at redshifts close to 0.003, it is possible for the O3] 1666
to be significantly affected by absorption from the Galactic Al2 1671 line
M101 has a redshift of $z\approx0.0008$ such that the O3] 1666.15 emission line
shifts to a line center of 1667.48 Å, but O3] can be further shifted towards
Al2 contamination by the rotational velocity of a given H2 region within M101.
These effects, combined with the low resolution of our spectra, cause the Al2
absorption to significantly affect the O3] profile.
The resulting effect is clearly seen in Figure <ref>:
the O3] line strengths are weak for most of the sample, where the blue Al2
absorption wing is clearly blended with the O3] emission.
Therefore, higher resolution spectra are needed to strongly detect the FUV O3]
nebular emission in M101 and other metal-rich targets.

§.§.§ The C/O Abundance Ratio

Determining relative abundances requires significant ($>3\sigma$) line detections.
While several of our C3] line detections meet this criteria, only Region 10 has
a significant O3] detection.
We, therefore, perform C/O abundance determinations only as an exploratory exercise,
noting that higher S/N and higher resolution spectra are needed for robust measurements.
Given the faint hints of O3] emission in Regions 9, 7, 6, and 4, significant
detections could likely be achieved with higher spectral resolution.
We note that we refrain from performing a 1-sigma upper limit
calculation because the contamination of the MW Al2 1671 absorption feature
would lead to an underestimation of the upper limit.=-2

To perform an exploratory C/O abundance analysis,
an initial approximation of C/O can be determined from
the C$^{+2}$/O$^{+2}$ ratios reported in Table <ref>.
However, the higher ionization potential of O$^{+2}$ relative to C$^{+2}$ (54.9 eV
versus 47.9 eV, respectively), means that an ionization correction factor is needed
to account for any unseen ionization states.
For the spectra shown in Figure <ref>, the C4 emission features all seem
to be well fit by a stellar wind P-Cygni profile, suggesting that we are not missing any
significant emission from C$^{+3}$ ions.
However, there could be contributions from lower ionization states.
Therefore, we determine C/O abundances using:

\begin{align}
{\frac{\mbox{C}}{\mbox{O} }} & = {\frac{\mbox{C}^{+2}}{\mbox{O}^{+2}}\ }\times \Bigg[{\frac{X(\mbox{C}^{+2})}{X(\mbox{O}^{+2})}}\Bigg]^{-1} \nonumber \\
& = {\frac{\mbox{C}^{+2}}{\mbox{O}^{+2}}\ }\times{\mbox{C ICF}},
\end{align}

where X(C$^{+2}$) and X(O$^{+2}$) are the C$^{+2}$ and O$^{+2}$ volume fractions, respectively.
Following the metallicity-dependent ICF method of Berg et al., 2019,
we ran cloudy 17.00 [Ferland et al., 2013] photoionization models to match the properties
spanned by M101.
We used  burst models for the input ionizing radiation field,
a range of burst ages ($10^{6.0}-10^{7.0}$ yrs),
ionization parameter ($-3.0 <$ log $U <-1.0$), and
metallicities ($Z = 0.004-0.02 = 0.2-1.0\ Z_{\odot}$),
with matched stellar and nebular metallicities ($Z_{\star} = Z_{neb}$).

Our cloudy models are shown in Figure <ref> comparing the [O3]
5007/[O2] 3727 emission line ratios, ionization parameters, and
ionization fractions of C$^{+2}$ and O$^{+2}$.
The top panel shows the increasing trend of [O3] 5007/[O2]
3727 with ionization parameter, which is a useful diagnostic.
The points are color-coded by model stellar/nebular metallicity, while the
shaded band denotes the corresponding models spanned by the range in burst age.
We fit each of the metallicity models with a polynomial of the shape:
log $U = c_3\cdot{x^2} + c_2\cdot{x} + c_1$, where
$x =$ log([O3] 5007/[O2] 3727).
The coefficients for these fits are listed in Table <ref>.
We note that the coefficients derived in this work are slightly different than
those reported in Berg et al., 2019 for common metallicities because we took the
average values of the models spanning different ages, while the previous work
fit the $t=10^{6.5}$ yr models only.
Using the observed [O3] 5007/[O2] 3727 ratios from
Croxall et al., 2016 and the metallicities reported in Berg et al., 2020, we determined
log $U$ values by interpolating between the polynomial fits defined for each
metallicity steps; the resulting values are reported in Table <ref>.

With log $U$ estimates in hand, we determined C ICFs for our sample.
As shown in the middle panel of Figure <ref>, both metallicity and
stellar age have moderate effects on the relative ionization fraction of C.
We fit each of the metallicity models with a polynomial of the shape:
C ICF = [X(C$^{+2}$)/X(O$^{+2}$)]$^{-1}$ =
$c_5\cdot{x^4} + c_4\cdot{x^3} + c_3\cdot{x^2} + c_2\cdot{x} + c_1$,
where $x =$ log $U$.
The coefficients for these fits are listed in Table <ref>.
We determined moderate C ICFs for our sample, showing a moderate
to small correction for our sample that lies between log$U=-2.75$ to $-2.30$.
We estimated the uncertainty in the ICF as the scatter amongst the different
models considered (relative abundances and burst age) at a given log$U$.
Ionization parameters, C ICFs, ionic C fractions, and corrected C/O ratios
are provided in Table <ref>.

C/O Emission-Line and Abundance Estimates

Property           Reg 4        Reg 5        Reg 6        Reg 7        Reg 8        Reg 9        Reg 10
0.144±0.129     0.078±0.253     0.163±0.162     0.073±0.177     0.192±0.328     0.198±0.233     0.183±0.014

E(B–V)                 0.185             0.093             0.042             0.180             0.070             0.049             0.231

$T_e$[O3] (K)    9540±90        9180±90         9440±70         9300±70         9580±80         10690±90        12790±170

C$^{+2}$/O$^{+2}$       0.21±0.68       0.36±0.50       0.18±0.24       0.48±0.44       0.16±0.44       0.17±0.28       0.23±0.10

log $O_{32}$            0.16±0.02       0.31±0.01       0.32±0.01       0.14±0.01       0.38±0.01       0.31±0.01       0.75±0.02

log $U$                 -2.70±0.09      -2.56±0.10      -2.58±0.09      -2.78±0.08      -2.55±0.08      -2.64±0.06      -2.23±0.06

C ICF                   0.791             0.862             0.851             0.750             0.862             0.824             0.981

C$^{+2}$/C ICF          0.714             0.768             0.785             0.734             0.781             0.766             0.830

log(C/O)   -0.79±0.63   -0.51±0.87   -0.81±0.36   -0.45±0.67   -0.87±0.58   -0.85±0.43    -0.65±0.16

log(N/O)                -1.23±0.19      -1.22±0.05      -1.12±0.03      -1.12±0.05      -1.09±0.06      -1.34±0.17      -1.33±0.11

12+log(O/H)             8.51±0.10       8.51±0.02       8.45±0.02       8.39±0.01       8.36±0.02       8.29±0.06     8.16±0.02

The top three rows give the reddening-corrected intensities of O3]
and C3] from the UV spectra relative to H$\beta$ from the optical
spectra, determined using the E(B–V) reddening reported in Croxall et al., 2016.
Note that the O3] flux is just the strong line of the doublet, the
1666 feature, but the C3] flux is the total of the 1907,1909 lines.
The bottom portion of the table lists the physical properties derived for the gas,
where optical-only properties are reported from Croxall et al., 2016 and Berg et al., 2020,
and UV-derived properties are from this work.
Only Region 10 has significant detections of both O3] and C3];
values for the other six regions are only estimates and are italicized to
denote this.

Coefficients for ICF Model Fits

1c 	    6c$Z(Z_{\odot})$

$y = f(x)$ 	      0.05     0.10    0.20     0.40     0.70     1.00
1llog $U$:

1l$x =$ log $O_{32}$

$c_3$ ..............      0.0798        0.0899        0.1045      0.1308      0.1592      0.1938

$c_2$ ..............      0.7253        0.7514  	    0.7932      0.8588      0.9462      1.0290

$c_1$ ..............      -3.026 	    -2.978	    -2.914      -2.914      -2.836      -2.801

1lC ICF:

1l$x =$ log $U$

$c_5$ ..............	    0.127	        0.111  	    0.098      0.081      0.070     0.056

$c_4$ ..............	    1.373   	    1.173	        1.001      0.815      0.671     0.528

$c_3$ ..............	    5.298   	    4.389 	    3.558      2.807      2.190     1.643

$c_2$ ..............	    9.070 	    7.267	        5.575      4.246      3.151     2.286

$c_1$ ..............	    7.023	        5.681	        4.415      3.520      2.792    	2.264

cloudy photoionization model fits of the form
The model grid and polynomial fits are shown in Figure <ref>.
The models are described in  <ref>.

The bottom panel of Figure <ref> shows the ionization fraction of
C$^{+2}$/(C$^+$+C$^{+2}$) vs O$^{+2}$/(O$^+$+O$^{+2}$).
If the singly- and doubly-ionized species are the dominant species of C and O,
then these ratios can be used directly as an independent ICF.
Using the O$^{+2}$/(O$^+$+O$^{+2}$) ratio derived from the optical spectra,
we predict the C$^{+2}$/(C$^+$+C$^{+2}$) ratios for our sample.
The C$^{+2}$/C ICF values are report in Table <ref>.
Compared to the our C ICFs derived from the ionization parameter,
the C$^{+2}$/C ICF values are always lower and less than one, which would
provide a more significant correction to our C/O abundances.

Note that the C and O abundances presented here have not been corrected
for the fraction of atoms embedded in dust.
Peimbert & Peimbert, 2010 have estimated that the depletion of O ranges between
roughly 0.08$-$0.12 dex, and has a positive correlation with O/H abundance.
C is also expected to be depleted in dust,
mainly in polycyclic aromatic hydrocarbons and graphite.
The estimates of the amount of C locked up in dust grains in the local
interstellar medium shows a relatively large variation depending on the
abundance determination methods applied <cit.>.
However, carbon and oxygen are often found to have similar amounts of depletion
onto dust grains, with corrections typically around 0.1 dex <cit.>.
Assuming the C and O depletions are similar, we take the gas-phase C/O ratio as
indicative of the total C/O abundance for this exercise.

Photoionization models used to derive the carbon ICF.
Models with $t=10^{6.5}$ yrs are plotted as squares and color-coded by their metallicity.
The best-fit polynomial function for each metallicity is plotted as a dashed line;
fit coefficients are reported in Table <ref>.
The range of models spanned by ages of $t=10^{6-7}$ yrs are designated
by the shaded bands, representing an estimate of the uncertainty.
Top: Ionization parameter, log $U$, as a function of the
[O3] 5007/[O2] 3727 emission line ratio.
Middle: Carbon ionization correction factor,
CICF = [X(C$^{+2}$)/X(O$^{+2}$)]$^{-1}$, as a function of log $U$.
Bottom: Another CICF, C$^{+2}$/(C$^+$+C$^{+2}$),
but derived as a function of the observed O$^{+2}$/(O$^+$+O$^{+2}$) ion fraction.
The observed line ratios and ion fractions from Croxall et al., 2016 and derived
log $U$ and ICF values are plotted as teal diamonds.

Carbon to oxygen ratio versus oxygen abundance for star-forming galaxies
and H2 regions.
Low-metallicity galaxies with CEL C/O measurements from the literature
<cit.> are plotted as gold diamonds.
Some additional CEL C/O abundances are plotted for individual H2 regions
from Garnett et al., 1999.
We also include RL C/O abundances measured for higher-metallicity H2 regions
from <cit.> and <cit.>.
The dashed line is the least-squares fit from Garnett et al., 1995,
while the  blue curve shows the empirical stellar curve from Nicholls et al., 2017.
Some of our preliminary CEL C/O measurements (turquoise diamonds) agree with the overall
increasing C/O–O/H trend, while others appear to be low-C/O outliers, but are
not significant enough to discern from the large scatter.

§.§.§ C/O Comparison to Previous Works

Several previous studies have analyzed the nebular abundances in M101.
Of these studies, four are particularly relevant to this work:
(1) Garnett et al., 1999 observed three H2 regions in M101 using the Faint
Object Spectrograph (FOS) on HST.
C3] 1907,1909 was detected significantly in all three,
but only upper limits could be placed on O3] 1661,1666.
As tabulated in Table <ref>, two of these regions overlap with our
H2 region sample in M101 (NGC5455, NGC5471).
(2) Esteban et al., 2009 obtained high-resolution optical spectroscopy with the
High Resolution Echelle Spectrometer <cit.> on the Keck I telescope
for two H2 regions in M101 and intermediate-resolution spectroscopy
of NGC 5447 using the Intermediate dispersion Spectrograph and Imaging System
(ISIS) spectrograph on the 4.2m WHT.
They measured significant C2 RLs in H1013 and NGC 5447b,
which overlaps with our UV CEL measurement.
(3) Skillman et al., 2020 measured significant C2 RLs in 10 H2 regions
in M101 using low-resolution optical spectroscopy taken with the MODS on the LBT.
(4) Esteban et al., 2020 used low-resolution optical spectroscopy taken with the
Optical System for Imaging and low-Intermediate-Resolution Integrated Spectroscopy
<cit.> spectrograph on the Gran Telescopio Canarias (GTC) to measure
the RLs in 11 H2 regions in M101, five of which have significant C2.
Of these regions, NGC 5455 and NGC 5471 are in common with Garnett et al., 1999,
Skillman et al., 2020, and this work, NGC 5447b is in common with Esteban et al., 2009,
Skillman et al., 2020, and this work, H1013 overlaps with the Skillman et al., 2020 sample,
while H37 has no comparison.
These previous works provide consistency checks on our measurements.

We conduct a preliminary comparison of our UV C/O abundance measurements to previous
studies in the C/O versus O/H plot in Figure <ref>.
In general, the study of the C/O versus O/H trend is of interest to improve our
understanding of C production and the conditions in which C emission is strongest
(e.g., high-redshift star-forming regions).
Theoretically, we expect C/O nucleosynthetic production to be constant via the
triple-$\alpha$ process.
However, observationally an increasing trend in C/O with O/H is clearly visible,
but with significant scatter.
This dispersion has prevented us from further discerning whether C/O increases
linearly or gradually with O/H (e.g., the dashed line from [Garnett et al., 1995]
or solid blue line from [Nicholls et al., 2017] in Figure <ref>) or has a
bimodal trend with a plateau at low O/H, similar to the N/O trend.

A further challenge to understanding the C/O–O/H relationship is the differing sources
of C/O measurements; at low-metallicity C/O is measured via the collisionally-excited
UV lines, while C/O is typically measured via the optical RLs at high-metallicity.
Therefore, our measurements, combined with the CEL measurements from Garnett et al., 1999,
provide an important extension of the CEL measurements at low-metallicity that would
allow an apples-to-apples assessment of the C/O–O/H trend if more significant
measurements were obtained.
Along these lines, our measurements agree well with the dispersion of the C/O trend in
Figure <ref>, but also lie at low C/O ratios on average.
In particular, the low-metallicity (gold diamonds) plateau has an average
log(C/O)$_{\rm CEL} = -0.71$, while the high-metallicity average
([Garnett et al., 1999] plus this work) log(C/O)$_{\rm CEL} = -0.54$.
Note that this value could be largely biased by the low average of our seven points
(log(C/O)$_{\rm CEL, B24} = -0.70$), whereas the Garnett et al., 1999 sample alone has
a much higher average (log(C/O)$_{\rm CEL, G99} = -0.36$).
In comparison, the high-metallicity average log(C/O)$_{\rm RL} = -0.29$, or roughly
0.1 dex higher than previous CEL C/O measurements, and much higher than the
values estimated here.
Despite the large uncertainties that remain, these results suggest that
(1) the increasing C/O trend is real and not just an artifact
of different measurement methods and
(2) that a C/O abundance discrepancy factor (ADF)
exists on the same order of magnitude, or larger, as the O/H ADF.
However, little more can be said given the large uncertainties on our C/O measurements.
Further progress to improve comparison of C/O from the UV CELs to those from optical RLs
and better constrain the source(s) of the ADF will require higher S/N and higher
spectral-resolution spectra.

Comparison of M101 CEL and RL C/O Measurements

2cName     R.A., Decl.     12+log(O/H)     log(C/O)      Method     Source
H1013$^1$      Reg 1     14:03:31.3, +54:21:05.8     8.57±0.02     -0.04±0.07     RL      Skillman et al., 2020

14:03:31.2, +54:21:14.8                     +0.22±0.13     RL      Esteban et al., 2009
H1052$^1$      Reg 2     14:03:34.1, +54:18:39.6     8.57±0.01     -0.08±0.04     RL      Skillman et al., 2020
NGC 5461       Reg 3     14:03:41.6, +54:19:08.5     8.48±0.02     -0.31±0.12     RL      Skillman et al., 2020

14:03:41.6, +54:19:04.4                     -0.03±0.13     CEL     Garnett et al., 1999$^2$
NGC 5462a      Reg 4     14:03:53.0, +54:22:06.8     8.45±0.05     -0.39±0.19     RL      Skillman et al., 2020

-0.79      CEL     This work
NGC 5462b      Reg 5     14:03:53.8, +54:22:10.8     8.43±0.01     -0.28±0.12     RL      Skillman et al., 2020

-0.51      CEL     This work
14:03:53.1, +54:22:06.4                     -0.12±0.10     RL      Esteban et al., 2020
NGC 5447       Reg 6     14:02:30.5, +54:16:09.9     8.42±0.01     -0.30±0.09     RL      Skillman et al., 2020

-0.81      CEL     This work
NGC 5455       Reg 7     14:03:01.1, +54:14:28.0     8.39±0.02     -0.34±0.11     RL      Skillman et al., 2020

-0.45      CEL     This work

14:03:01.2, +54:14:27.0                     -0.26±0.14     CEL     Garnett et al., 1999$^2$

14:03:01.2, +54:14:29.4                     -0.22±0.11     RL      Esteban et al., 2020
NGC 5447       Reg 8     14:02:27.8, +54:16:25.3     8.35±0.01     -0.25±0.17     RL      Skillman et al., 2020

-0.87      CEL     This work

14:02:28.7, +54:16:25                       -0.15±0.12     RL      Esteban et al., 2009

14:02:28.1, +54:16:26.9                     -0.27±0.11     RL      Esteban et al., 2020
H37$^1$                  14:02:17.7, +54:22:34.1     8.30±0.01     -0.14±0.14     RL      Esteban et al., 2020
H1216$^1$      Reg 9     14:04:10.9, +54:25:19.2     8.26±0.03     -0.63±0.50     RL      Skillman et al., 2020

-0.85      CEL     This work
NGC 5471       Reg 10    14:04:29.0, +54:23:48.6     8.14±0.03     -0.47±0.46     RL      Skillman et al., 2020

-0.65±0.16     CEL     This work

14:04:29.0, +54:23:48.8                     -0.67±0.05     CEL     Garnett et al., 1999$^2$

14:04:29.0, +54:23:49.0                     -0.20±0.11     RL      Esteban et al., 2020

Comparison of M101 CEL C/O abundances to RL C/O abundances from the literature and this work.
Column 1 lists the common name of the H2 region in M101, along with the
region name of this work, while coordinate R.A. and Decl. (J2000) is given in Column 2.
Column 3 lists the measured 12$+$O/H abundances from Croxall et al., 2016,
while Column 4 lists the C/O abundance measured using the method in Column 5 by the reference in Column 6.
Note that estimates of log(C/O) using insignificant detections on O3] in this work are
identified with italicized values.

$^1$Name adopted from the catalog of Hodge et al., 1990.
$^2$Note that the values from Garnett et al., 1999 are those calculated with an R$_V$ value of 3.1.
Garnett et al., 1999 also calculated values assuming an R$_V$ value of 5, which can give significantly
different results.

§ CONCLUSIONS

We present low-resolution HST/COS FUV spectroscopy of 9 H2 regions in
the star-forming disk of spiral galaxy M101.
These regions were selected for their high surface brightnesses and existing direct
abundances and C2 4267 RL detections from the optical CHAOS survey,
with the goal of detecting the FUV O3] 1666 and C3]
1907,1909 CELs.
Strong C3] emission lines were detected for the 7 lowest-metallicity H2
regions in our sample (12+log(O/H)$\lesssim8.5$), but O3] 1666 was significantly
detected in only the most metal-poor region (Reg 10; 12+log(O/H) = 8.16).
The other six H2 regions have tentative O3] detections that would be
improved with higher spectral resolution due to the narrow intrinsic widths of the lines
and the contamination of the Milky Way Al2 1671 ISM absorption line.

We performed an exploratory C/O abundance analysis with the limited O3] and
C3] line detections.
To convert the C/O emission line ratios to relative abundances, we produced new analytic
functions of the carbon ICF as a function of metallicity that are calibrated to
the conditions of the M101 H2 regions.
To do so, we created a custom grid of  photoionization models
using  single-burst models as inputs.
With our tentative C/O abundance measurements, we compared to the existing sample of C/O
measurements for M101, consisting of 17 optical RL measurements and 3 FUV CEL measurements for
11 different H2 regions (see Table <ref>).
Our preliminary C/O measurements add 7 CEL points to this sample, but five of these
points our low C/O outliers from the general C/O–O/H trend (Figure <ref>).
While the FUV C/O abundances can provide useful constraints on the chemical
enrichment of galaxies, higher-S/N and higher-resolution CEL measurements are clearly
needed to robustly constrain the high-metallicity end of the C/O–O/H and
compare to RL C/O abundances.

Despite the lack of strong O3] CELs, the resulting FUV spectra revealed
numerous continuum features, such as the massive-star wind features, that are useful
for characterizing the ionizing stellar population.
We fit stellar continuum with both  and
luminosity-weighted stellar burst models to determine the characteristic
age and metallicity of the massive star populations, finding
a range of ages (2–10 Myr) and metallicities (0.5–1.0 $Z_\odot$) that
seem to correlate with the He and N emission and P-Cygni profiles observed.
In particular, the youngest stellar populations are strongly correlated with
broad Wolf-Rayet (WR) He2 1640 emission and enhanced N/O abundances
(log(N/O)$>-1.2$), as determined from the optical
[N2] 6584/[O2] 3727 ratio.
Detections of WR stars are further supported by visual confirmation of the
WR “blue bump" feature around 4650 in the optical spectra, where the strong
N2 4634,4641,4643 emission is indicative of nitrogen-type WR (WN) stars.
As a result, this special phase of stellar evolution may be responsible for the
N4 1483,1486 emission and the excess emission in the N4 1718
stellar wind feature, which in turn biases the  and
stellar continuum fits to higher metallicities relative to the gas-phase metallicities.
Therefore, the WR H2 regions of M101 presented here provide a strong test-bed
to constrain future WR atmosphere and evolution models, and their incorporation to
stellar population models.

Finally, we observed additional P-Cygni and emission features in the $\sim900-1200$ Åregime of our M101 spectra that further support strong contributions from WN stars
(e.g., S6 933,945, N4 955, N3 991, O6 1032,1038,
P5 1118,1128, and C3 1175), and note a lack of strong
C3 and C4 features in this regime (e.g., C3 977,
C4 1197, C4 1135, C3 1140, C4 1169)
that would be expected if WC stars were present.
Therefore, we recommend the $\sim 900-1200$ Å stellar continuum as a powerful
diagnostic regime for identifying the presence of WR stars in the integrated spectra
of star-forming regions and galaxies.

HST (COS), LBT (MODS)

[Astropy Collaboration et al., 2013, Astropy Collaboration et al., 2018, Astropy Collaboration et al., 2022],
[Green, 2018],
[Kluyver et al., 2016],
[Kelly, 2007],
[Markwardt, 2009],
MODS reduction Pipeline,
[Bradley et al., 2021],
version 1.2.2 [Newville et al., 2018],
version 1.1.14 [Luridiana et al., 2015],
version 2.2 [Stanway & Eldridge, 2018],
[Leitherer et al., 2014],
version 1.26 [Harris et al., 2020]

The authors are grateful to the referee for insightful comments and questions that
led to better parameter characterization and an overall improved paper.
DAB is grateful for the support for this program, HST-GO-15126, that was provided by
NASA through a grant from the Space Telescope Science Institute,
which is operated by the Associations of Universities for Research in Astronomy,
Incorporated, under NASA contract NAS5-26555.

This work was based in part on observations made with the Large Binocular Telescope (LBT). The LBT is an international collaboration among institutions in the United States, Italy, and Germany. The LBT Corporation partners are: the University of Arizona on behalf of the Arizona university system; the Istituto Nazionale di Astrofisica, Italy; the LBT Beteiligungsgesellschaft, Germany, representing the
Max Planck Society, the Astrophysical Institute Potsdam, and
Heidelberg University; the Ohio State University; and the
Research Corporation, on behalf of the University of Notre
Dame, the University of Minnesota, and the University of
Virginia. This research used the facilities of the Italian Center for Astronomical Archive (IA2) operated by INAF at the Astronomical Observatory of Trieste.

This paper made use of the modsIDL spectral data reduction reduction pipeline
developed in part with funds provided by NSF Grant AST-1108693 and a generous
gift from OSU Astronomy alumnus David G. Price through the Price Fellowship in
Astronomical Instrumentation.

All the HST data used in this paper can be found in MAST:
[ 10.17909/qv80-ph56]http://dx.doi.org/ 10.17909/qv80-ph56.

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