# X-ray Emission of Nearby Low-mass and Sun-like Stars with Directly Imageable
Habitable Zones

Breanna A. Binder Department of Physics and Astronomy, California State
Polytechnic University, Pomona, CA, USA  Sarah Peacock University of Maryland,
Baltimore County, Baltimore, MD, 21250, USA  NASA Goddard Space Flight Center,
Greenbelt, MD 20771, USA  Edward W. Schwieterman Department of Earth and
Planetary Sciences, University of California, Riverside, CA, USA  Blue Marble
Space Institute of Science, Seattle, WA, USA  Margaret C. Turnbull SETI
Institute, Carl Sagan Center for the Study of Life in the Universe, Off-Site:
Madison, WI 53713, USA  Azariel Y. Virgen  Department of Physics and
Astronomy, California State Polytechnic University, Pomona, CA, USA  Stephen
R. Kane Department of Earth and Planetary Sciences, University of California,
Riverside, CA, USA  Alison Farrish NASA Goddard Space Flight Center,
Greenbelt, MD 20771, USA  Katherine Garcia-Sage Heliophysics Division, NASA
Goddard Space Flight Center, Greenbelt, MD, USA

###### Abstract

Stellar X-ray and UV radiation can significantly affect the survival,
composition, and long-term evolution of the atmospheres of planets in or near
their host star’s habitable zone (HZ). Especially interesting are planetary
systems in the solar neighborhood that may host temperate and potentially
habitable surface conditions, which may be analyzed by future ground and
space-based direct-imaging surveys for signatures of habitability and life. To
advance our understanding of the radiation environment in these systems, we
leverage $\sim$3 Msec of XMM-Newton and Chandra observations in order to
measure three fundamental stellar properties at X-ray energies for 57 nearby
FGKM stellar systems: the shape of the stellar X-ray spectrum, the luminosity,
and the timescales over which the stars vary (e.g., due to flares). These
systems possess HZs that will be directly imageable to next-generation
telescopes such as the Habitable Worlds Observatory and ground-based Extremely
Large Telescopes (ELTs). We identify 29 stellar systems with $L_{\rm
X}$/$L_{\rm bol}$ ratios similar to (or less than) that of the Sun; any
potential planets in the habitable zones of these stars therefore reside in
present day X-ray radiation environments similar to (or less hostile than)
modern Earth, though a broader set of these targets could host habitable
planets. An additional 19 stellar systems have been observed with the Swift
X-ray Telescope; in total, only $\sim$30% of potential direct imaging target
stars has been observed with XMM-Newton, Chandra, or Swift. The data products
from this work (X-ray light curves and spectra) are available via a public
Zenodo repository (doi: 10.5281/zenodo.11490574).

Planet hosting stars (1242); Stellar X-ray flares (1637); F dwarf stars (516);
G dwarf stars (556); K dwarf stars (876); M dwarf stars (982); Habitable
planets(695)

††facilities: CXO, XMM, Swift (XRT), Exoplanet Archive††software: astropy
(Astropy Collaboration et al., 2013, 2018, 2022)

##  1 Introduction

In the coming decades, ground- and space-based direct imaging will afford the
best potential for detecting and spectrally characterizing terrestrial planets
in the habitable zones of our nearest neighbors (Turnbull et al., 2021;
Kopparapu et al., 2018; Li et al., 2021; Kane et al., 2018; Currie et al.,
2023). While the James Webb Space Telescope (JWST) has and will continue to
characterize some rocky exoplanets (e.g., Greene et al., 2023; Zieba et al.,
2023; Lincowski et al., 2023; Lustig-Yaeger et al., 2023; Moran et al., 2023),
its ability to spectrally examine temperate terrestrial—and therefore
potentially habitable—exoplanets is limited primarily to target systems with
transiting planets and late M dwarf host stars. This is due to the lack of
coronagraphic instrumentation on JWST to image rocky planets around more
massive, Sun-like stars, and because the characterization of terrestrial
planets through transit observations is most favorable for M dwarf hosts,
which have a combination of favorable transit probabilities, planet-star size
ratios, and planet-star separations that allow for the probing of low
atmospheric altitudes when accounting for the effects of refraction (García
Muñoz et al., 2012; Bétrémieux & Kaltenegger, 2013, 2014). Moreover, some key
candidate biosignature molecules like O2 cannot be detected directly at Earth-
like abundances with JWST via transit spectroscopy even for the most favorable
targets (e.g., Lustig-Yaeger et al., 2019; Pidhorodetska et al., 2020; Meadows
et al., 2023), but could be more easily detected by reflected light
observations of planets in nearby systems (The LUVOIR Team, 2019; Gaudi et
al., 2020).

The occurrence rates derived from Kepler indicate that terrestrial planets are
likely present within the habitable zones (HZs) of a large fraction of nearby
stars (e.g., Bryson et al., 2021; Kunimoto & Matthews, 2020; Dattilo et al.,
2023; Fulton et al., 2017; Burke et al., 2015; Kane et al., 2014). The
Astro2020 Decadal Survey has identified the discovery and characterization of
these worlds as the highest priority for NASA’s investments (National
Academies of Sciences & Medicine, 2023), and the Habitable Worlds
Observatory111https://habitableworldsobservatory.org/home will build on the
LUVOIR and Habitable Exoplanets Observatory concepts (The LUVOIR Team, 2019;
Gaudi et al., 2020) to realize this goal. Due to detection biases, the vast
majority of nearby temperate rocky planets amenable for direct imaging have
not yet been discovered, but will be revealed by HWO in its survey phase
(e.g., The LUVOIR Team, 2019; Gaudi et al., 2020), precursor science with
extreme precision radial velocity (EPRV) measurements (Crass et al., 2021;
Morgan et al., 2021), or direct imaging with upcoming 30-m class telescopes
(Fujii et al., 2018; Currie et al., 2023).

A detailed understanding of the target stars themselves is a strong
prerequisite for maximizing the science yield of exoplanet imaging missions,
including extracting the statistical features of the overall planet population
(e.g., Van Eylen et al., 2014) as well as characterizing the climates,
atmospheric composition and stability (Linsky & Redfield, 2024; Louca et al.,
2023), and potential biosignatures of individual planets (Rugheimer et al.,
2015; Schwieterman et al., 2018). Evaluating planetary habitability,
specifically, must consider myriad factors beyond circumstellar location
including the planet’s bulk composition (Hinkel & Unterborn, 2018; Marounina &
Rogers, 2020; Unterborn et al., 2023) and stellar X-ray and extreme
ultraviolet (XUV; 1-912 Å; corresponding to an energy range of 0.01-12 keV)
emission (France et al., 2016; Garcia-Sage et al., 2017; Luger & Barnes, 2015;
Richey-Yowell et al., 2019). Planets orbiting in or near their star’s HZ can
experience high levels of XUV radiation, which directly influences atmospheric
survival, composition, and long-term evolution (Peacock et al., 2020;
Johnstone et al., 2021). In some cases, planetary atmospheres and surface
volatiles can be entirely driven away by XUV irradiation (e.g., Sanz-Forcada
et al., 2011; Fromont et al., 2024; Luger & Barnes, 2015). Further, the
atmospheric states of planets orbiting M dwarf stars may be profoundly
different than for those orbiting Sun-like FGK stars due their to extended
super-luminous pre-main sequence phase (Baraffe et al., 2015) during which XUV
radiation could ablate away planetary atmospheres or lead to atmospheric
oxidation via massive H escape (Garcia-Sage et al., 2017; Luger & Barnes,
2015; Krissansen-Totton & Fortney, 2022). Outgassing could replenish these
atmospheres (Kite & Barnett, 2020; Swain et al., 2021; Unterborn et al., 2022;
Kane et al., 2020a) though the retention and photochemistry of the atmosphere
would continue to be controlled by the host star spectrum and activity
(Roettenbacher & Kane, 2017; France et al., 2020, 2016; Loyd et al., 2018a, b;
Do Amaral et al., 2022).

While X-ray detections have been made for some HWO target stars (Harada et
al., 2024) and more detailed studies of the X-ray properties of low-mass
nearby stars have been performed (Brown et al., 2023; France et al., 2020;
Loyd et al., 2018b), no comprehensive and uniform X-ray analysis of potential
direct imaging target stars has yet been done. In this paper, we present a
comprehensive analysis of archival X-ray observations for all FGKM stellar
systems with HZs that will be directly-imageable with upcoming telescopes such
as HWO and ground-based Extremely Large Telescopes (ELTs). A follow-up paper
yielding panchromatic XUV–IR spectra (including modeled EUV wavelengths) for
each one of these systems is forthcoming (Peacock et al., in prep). We
describe our sample and the archival data in Sections 2 and 3. Our analysis
procedure is detailed in Section 4. We extract light curves for stars detected
at high significance and identify periods of count rate variability (such as
flaring events or epochs of elevated count rates). We then extract X-ray
spectra for stars, separating variable- from non-variable epochs, and fit the
resulting spectra with one-, two-, or three-component thermal plasma models.
Finally, the average best-fit spectral models are used to convert observed
count rates (or count rate upper limits, in the case of non-X-ray-detected
stars) to luminosities. In Section 5, we discuss the evolution of the $L_{\rm
X}$/$L_{\rm bol}$ ratio with effective temperature and stellar age, and
present an extensive appendix describing the X-ray observations and a summary
of the known stellar physical parameters and exoplanet systems from the
literature for each stellar system in our sample.

##  2 The Sample

To identify the highest priority targets for “deep dive” investigations of
nearby stars and their planetary systems, we have used the predicted imaging
and spectroscopy capability of (1) NASA’s next generation space-based Great
Observatories such as HWO, and (2) ground-based ELTs. To first order, the
performance of these facilities is described in terms of two specifications:
the inner working angle (IWA) — the smallest planet-star angular separation
that can be imaged (expressed in milliarcseconds, or mas), and the minimum
planet to star flux ratio (or fractional planet brightness, FPB) — where
starlight suppression technology and postprocessing algorithms can separate
the planet signal from stellar photons and systematic noise. Intriguingly (and
somewhat vexingly), for HZ planets the IWA requirement favors stars that are
more luminous and therefore have wider HZs, while the limiting flux ratio
favors target stars that are less luminous and therefore outshine their HZ
planets by a smaller factor (Turnbull et al., 2012). A third specification,
planet limiting magnitude ($V_{\rm lim}$), determines which of these targets
planets would be bright enough to obtain spectra in a reasonable amount of
observing time – and this requirement simply favors the nearest stars.

To form our preliminary target list, for space-based missions we chose stars
for which the HZ (0.95 - 1.7 AU scaled by $\sqrt{L_{\rm bol}}$) falls at least
partially outside an inner working angle of 58 mas, and for which Earth-sized
planets of albedo $\sim$0.3 would be brighter than $V_{\rm lim}<30$, with a
fractional planet brightness (FPB) $>4\times 10^{-11}$. For ground-based
observatories, we examined the 25 stars within 10 pc having the most favorable
FPB whose HZs are at least partially visible outside an IWA of $\sim$30 mas,
consistent with the literature (Brandl et al., 2018; Quanz et al., 2015). This
results in a well-defined list of 175 target stars whose HZs fall at least
partially within the detection space of upcoming exoplanet direct imaging
observatories. We compared our target list with the potential target stars for
HWO (Mamajek & Stapelfeldt, 2024), who used slightly different cut-off
requirements for $V_{\rm lim}$, FPB and IWA. Our list contains more M-dwarfs
(more favorable targets for ELTs) while the Mamajek & Stapelfeldt (2024) list
includes more F stars (more favorable for space-based observations). Mamajek &
Stapelfeldt (2024) also excludes stars in close binaries or higher-order
multiple systems, while we produce a more inclusive X-ray catalog of nearby
stars, assuming technical challenges posed by binaries may eventually be
surmounted. There are 229 unique stars when combining the two lists, which we
summarize in Table 1.

Table 1: Potential Target Stars for Future HZ Imaging Surveys

HD | Common | Distance | Spectral | From | Potential
---|---|---|---|---|---
ID | Name | (pc) | Type | This Work? | HWO Targetaafootnotemark:
166 | V439 And | 13.77 | G8V | ✓ | ✓
693 | 6 Cet | 18.89 | F8V |  | ✓
739 |  $\theta$ Scl | 21.72 | F5V |  | ✓
1326 | GX And | 3.57 | M1.5V | ✓ |
1581 |  $\zeta$ Tuc | 8.61 | F9.5V | ✓ | ✓
⋮ | ⋮ | ⋮ | ⋮ | ⋮ | ⋮

Note. — aFrom Mamajek & Stapelfeldt (2024). Select entries are shown to
illustrate the table form and content. The full machine-readable table (229
rows) is available online from the journal.

Using this input list of 229 viable targets for direct imaging campaigns, we
searched the archives of the XMM-Newton Observatory and the Chandra X-ray
Observatory for publicly available data. Many M stars (and some K stars) in
this initial target list were previously studied as part of the Measurements
of the Ultraviolet Spectral Characteristics of Low-mass Exoplanetary systems
(MUSCLES) and Mega-MUSCLES Hubble Space Telescope Treasury programs, which
characterized the X-ray/UV emission from K and M exoplanet host stars. We
incorporate these results, presented in Brown et al. (2023), into our work
here (see Section 5). We retrieve 192 XMM-Newton and Chandra observations for
57 stellar systems containing Sun-like and low-mass stars from the X-ray
archives, amounting to $\sim$3 Msec of observing time (not including the
MUSCLES/Mega-MUSCLES observations). An additional 19 stellar systems have been
observed by Swift. Figure 1 shows a map of the 3D spatial distribution of
these nearby stars and indicates which stars have X-ray observations
available.

There are three M-dwarfs highlighted in Mamajek & Stapelfeldt (2024) as strong
potential HWO targets (Lacaille 8760, Lacaille 9352, and Lalande 21185). Both
Lacaille 8760 and Lacaille 9352 were observed by Swift and are included in
this work. A detailed study of the panchromatic spectrum and habitability
potential of Lalande 21185 is in preparation (Schwieterman et al., in prep)
and is not included in the sample here. Roughly half of the target stars
contained in our study are F- and G-type stars. Table 3.1 provides a summary
of the stellar physical parameters for all stars included in our study,
including the effective temperature ($T_{\rm eff}$), spectral type (from Gray
et al., 2006), stellar mass ($M$), bolometric luminosity ($L_{\rm bol}$), and
approximate age (in Gyr). All data were compiled using the NASA Exoplanet
Archive222See https://exoplanetarchive.ipac.caltech.edu/. Right ascension
(R.A.), declination (Dec.), and distances were taken from Gaia DR3 (Gaia
Collaboration et al., 2023). More detailed discussion of each stellar system
is provided in Appendix A.

Figure 1: Map of nearby potential direct imaging target stars, color-coded by
spectral type. Stars with blue halos have available X-ray observations. The
yellow star at the center of the diagram represents the Sun, with white
concentric rings showing distances of 5 pc, 10 pc, and 15 pc. An animated
version of this figure (8 seconds in duration) is available from the online
journal showing the 3D distribution through different elevation and azimuthal
angles (colors and symbols are the same as the static figure).

##  3 X-ray Data

We retrieved all publicly available imaging observations for our target stars
from the XMM-Newton333See https://nxsa.esac.esa.int/nxsa-web/#home and
Chandra444See https://cda.harvard.edu/chaser/ archives. Table 3 summarizes the
number of observations and the approximate exposure times utilized in this
work. Below, we provide a more detail description of how the observations were
processed and the methods used to extract light curves and spectra for each
facility.

###  3.1 XMM-Newton

We retrieved XMM-Newton observations from the XMM-Newton Science Archive and
reprocessed the data from the raw evt1 files using SAS
18555https://www.cosmos.esa.int/web/xmm-newton/sas. The European Photon
Imaging Camera (EPIC) on board XMM-Newton carries a set of three X-ray CCD
arrays: one array uses pn-CCDs (referred to as PN) and two utilize metal oxide
semi-conductor CCDs (referred to as MOS1 and MOS2). The PN data were
reprocessed with epproc and the MOS data were reprocessed using emproc and
filtered on an energy range of 0.2-15 keV. Due to the brightness of many of
the stars, pile-up (when two or more photons fall on a pixel in less than the
readout time) is a significant concern. The PATTERN parameter records the
number and pattern of CCD pixels that are triggered for a given event; we use
the recommended PATTERN$=$0 to mitigate pile-up effects666See the XMM-Newton
ABC Guide at https://heasarc.gsfc.nasa.gov/docs/xmm/sl/epic/image/sas_cl.html.
Background light curves were extracted using evselect and inspected for
background flares, and good time intervals (GTIs) were generated to reject
observation times exhibiting strong background flaring events.

We inspected the reprocessed, cleaned images for evidence of an X-ray source
at the location of the target star. If the target star was detected at
$\gtrsim$500 counts in the PN image, we extracted source light curves (binned
to 100 s) to search for count-rate variability. We used evselect to extract
both PN and MOS1/2 spectra for stars with $\gtrsim$2000 net counts, rmfgen to
generate redistribution matrix files (RMFs), and arfgen to generate ancillary
response functions (ARFs) for all available observations. All spectra were
binned to contain at least 25 counts per bin.

The majority ($\sim$83%) of the observations analyzed in this work are from
the XMM-Newton archive, containing $\sim$75% of the total exposure time.
However, with a spatial resolution of $\sim$15′′, only stellar binaries with
angular separations larger than $\sim$30′′ can be resolved with the PN camera.
Figure 2 (left panel) shows a 10 ks XMM-Newton/PN image of the binary system
GL 412, where the two components (and a third, unidentified X-ray source) are
resolved.

Table 2: Properties of Sample Stars

Star | Alternate | R.A. | Dec. | Spectral | Distance | Mass | $T_{\rm eff}$ | $L_{\rm bol}$ | Age∗ |
---|---|---|---|---|---|---|---|---|---|---
Name | Name(s) | (J2000) | Type∗ | (pc) | ($M_{\odot}$) | (K) | ($L_{\odot}$) | (Gyr) | References
$\eta$ Crv | HIP 61174 | 12:32:03.75 | -16:11:46.55 | F2V | 17.96$\pm$0.17 | 1.48${}^{+0.29}_{-0.19}$ | 6854${}^{+124}_{-99}$ | 4.62${}^{+0.22}_{-0.25}$ | 1.5${}^{+0.2}_{-0.4}$ | 1,2,50
$\xi$ Oph | HIP 84893 | 17:21:00.68 | -21:06:49.81 | F2V | 17.06$\pm$0.10 | 1.44${}^{+0.30}_{-0.20}$ | 6756${}^{+110}_{-127}$ | 4.13${}^{+0.17}_{-0.18}$ |  $\cdots$ | 1,2
$\upsilon$ And A | HD 9826 | 01:36:47.60 | 41:24:13.54 | F8V | 13.41$\pm$0.06 | 1.29$\pm$0.04 | 6157$\pm$112 | 3.35${}^{+0.10}_{-0.14}$ | 3.12 | 1,2,29
$\iota$ Hor | HR 810, HD 17051 | 02:42:34.03 | -50:47:57.54 | F9V | 17.32$\pm$0.02 | 1.17${}^{+0.17}_{-0.15}$ | 6146${}^{+106}_{-184}$ | 1.74${}^{+0.04}_{-0.06}$ | 0.47; 2.72 | 1-4,9
$\nu$ Phe | HIP 5862 | 01:15:12.13 | -45:31:51.14 | F9V | 15.17$\pm$0.04 | 1.15${}^{+0.18}_{-0.14}$ | 6119${}^{+130}_{-154}$ | 2.00$\pm$0.06 | 2.88 | 1,2,3,9
$\gamma$ Pav | HIP 105858 | 21:26:26.81 | -65:21:45.50 | F9V | 9.27$\pm$0.04 | 1.15${}^{+0.18}_{-0.14}$ | 6109${}^{+110}_{-176}$ | 1.47$\pm$0.05 | 1,7.25 | 1,3,44,45
$\beta$ Vir | HIP 57757 | 11:50:42.51 | 01:45:48.67 | F9V | 11.12${}^{+0.07}_{-0.06}$ | 1.13${}^{+0.16}_{-0.14}$ | 6071${}^{+120}_{-90}$ | 3.77${}^{+0.11}_{-0.14}$ | 2.96 | 1,2
GJ 1095 | HIP 35136 | 07:15:50.19 | +47:14:20.89 | F9V | 16.86${}^{+0.19}_{-0.20}$ | 1.12 | 5866 | 1.54$\pm$0.06 | 1, 7 | 1,2,10
LHS 237 | GJ 288, HIP 37853 | 07:45:34.63 | -34:09:53.86 | F9V | 15.20$\pm$0.13 | 1.11 | 5744 | 1.51$\pm$0.06 |  $\cdots$ | 1,2
$\iota$ Per | HIP 14632 | 03:09:06.10 | +49:36:46.35 | F9.5V | 10.51$\pm$0.03 | 1.08${}^{+0.18}_{-0.10}$ | 5938${}^{+141}_{-169}$ | 2.23${}^{+0.04}_{-0.10}$ | 4.20$\pm$0.48 | 1,2,9
$\beta$ Hyi | HIP 2021 | 00:25:55.80 | -77:15:10.06 | G0V | 7.47$\pm$0.03 | 1.14${}^{+0.01}_{-0.02}$ | 5839 | 3.70$\pm$0.05 | 6.32 | 1,2,9
$\beta$ Com | HIP 64394 | 13:11:51.43 | +27:52:55.58 | G0V | 9.18$\pm$0.03 | 1.10${}^{+0.14}_{-0.13}$ | 5969${}^{+188}_{-100}$ | 1.41${}^{+0.05}_{-0.04}$ |  $<$1.12, 1.7 | 1,2,9
LHS 208 | HD 39091, $\pi$ Men | 05:37:11.89 | -80:27:52.05 | G0V | 18.27$\pm$0.02 | 1.094$\pm$0.039 | 6037$\pm$45 | 1.44$\pm$0.02 | 2.98${}^{+1.40}_{-1.30}$ | 5
$\rho$ CrB | GJ 9537, HIP 78459 | 16:01:02.41 | +33:18:00.28 | G0V | 17.47$\pm$0.02 | 1.05${}^{+0.12}_{-0.15}$ | 5833${}^{+141}_{-124}$ | 1.81${}^{+0.06}_{-0.04}$ | 11.04 | 1,2,9
GL 672 | HIP 84862 | 17:20:39.74 | +32:27:47.22 | G0V | 14.54$\pm$0.02 | 1.05${}^{+0.16}_{-0.10}$ | 5817${}^{+120}_{-84}$ | 1.29${}^{+0.04}_{-0.03}$ | 12.04 | 1,2,9
$\chi^{1}$ Ori | HIP 27913, GJ 222 | 05:54:22.76 | +20:16:33.07 | G0V | 8.66$\pm$0.08 | 1.028${}^{+0.03}_{-0.028}$ | 5822 | 1.17$\pm$0.06 | 0.35; 4.32 | 1,2,6,9
GL 788 | HIP 100017 | 20:17:32.60 | +66:51:18.04 | G0V | 17.57$\pm$0.52 | 1.02$\pm$0.03 | 5886 | 1.18$\pm$0.10 | 3.76${}^{+1.88}_{-1.92}$ | 2,9
47 UMa | GJ 407, HD 95128 | 10:59:27.53 | +40:25:49.80 | G0V | 13.80$\pm$0.03 | 1.005$\pm$0.047 | 5829$\pm$95 | 1.58$\pm$0.04 | 6.48${}^{+1.44}_{-1.04}$ | 9,29
GL 620.1A | HD 147513 | 16:24:01.39 | -39:11:34.68 | G1V | 12.903$\pm$0.02 | 1.06${}^{+0.12}_{-0.14}$ | 5873${}^{+110}_{-103}$ | 1.032${}^{+0.003}_{-0.004}$ | 2; $<$1; 0.4 | 1-4,9,19
GL 311 |  $\pi^{1}$ UMa, HD 72905 | 08:39:11.64 | +65:01:16.67 | G1.5V | 14.27${}^{+0.14}_{-0.15}$ | 1.01 | 5920 | 1.03$\pm$0.05 | 0.2 | 2,6,8
$\alpha$ Cen A | GJ 559A, LHS 50 | 14:39:36.49 | -60:50:02.37 | G2V | 1.347$\pm$0.003 | 1.06${}^{+0.05}_{-0.04}$ | 5801 | 1.61$\pm$0.07 | 7.84${}^{+1.08}_{-1.28}$ | 2,9
44 Boo A |  $\iota$ Boo, HIP 73695 A | 15:03:46.74 | +47:39:15.89 | G2V | 12.8$\pm$0.2 | 1.04$\pm$0.1 | 5877 | 1.552 | 1.5 | 46,47,48
18 Sco | GJ 616, HIP 79672 | 16:15:37.52 | -08:22:17.91 | G2V | 14.13$\pm$0.02 | 1.04${}^{+0.12}_{-0.14}$ | 5791${}^{+117}_{-93}$ | 1.09${}^{+0.04}_{-0.02}$ | 5.84${}^{+1.88}_{-1.96}$ | 1,9,10
51 Peg | HIP 113357, GJ 882 | 22:57:28.22 | +20:46:08.78 | G2V | 15.46$\pm$0.03 | 1.03${}^{+0.17}_{-0.09}$ | 5758${}^{+102}_{-120}$ | 1.37${}^{+0.02}_{-0.05}$ | 6.76 | 1,2
$\zeta^{1}$ Ret | LHS 171, HIP 15330 | 03:17:49.26 | -62:34:20.76 | G2V | 12.12$\pm$0.08 | 0.96${}^{+0.03}_{-0.06}$ | 5699 | 0.80$\pm$0.05 | 1.56$\pm$0.44 | 2,9
HD 136352 | GJ 582, $\nu^{2}$ Lupi | 15:21:45.55 | -48:19:07.88 | G2V | 14.68$\pm 0.02$ | 0.87$\pm$0.04 | 5664$\pm$61 | 1.04$\pm$0.06 | 12.3${}^{+1.2}_{-2.9}$ | 11,12
GL 327 | HIP 43726 | 08:54:17.50 | -05:26:03.56 | G3V | 17.42${}^{+0.44}_{-0.47}$ | 1.09$\pm$0.04 | 5790 | 1.09${}^{+0.05}_{-0.04}$ | 1.32 | 2,9
$\mu$ Ara | HD 160691 | 17:44:08.68 | -51:50:05.65 | G3IV/V | 15.60$\pm$0.03 | 1.03${}^{+0.15}_{-0.11}$ | 5772${}^{+119}_{-84}$ | 1.90${}^{+0.04}_{-0.06}$ | 5.7$\pm$0.6 | 1,2,7
$\kappa^{1}$ Cet | HD 20630 | 03:19:21.98 | +03:22:14.22 | G5V | 9.14$\pm$0.03 | 1.02${}^{+0.16}_{-0.10}$ | 5712${}^{+161}_{-124}$ | 0.86${}^{+0.02}_{-0.03}$ | 0.35; 0.7; 2.2 | 1,2,6,20,21
HD 140901 | GL 599 | 15:47:28.54 | -37:55:02.15 | G7IV/V | 15.25$\pm$0.01 | 1.06$\pm$0.05 | 5587${}^{+153}_{-125}$ | 1.22$\pm$0.03 | 3.2$\pm$2.8 | 1,49
GJ 777A | LHS 3510, HD 190360 | 20:03:38.25 | +29:53:40.08 | G7IV/V | 16.01$\pm$0.02 | 0.99$\pm$0.08 | 5552$\pm$11 | 1.20$\pm$0.06 | 13.4 | 1,2,9
$\xi$ Boo A | 37 Boo, HD 131156 | 14:51:23.53 | +19:06:00.97 | G7V | 6.71$\pm$0.16 | 0.90$\pm$0.04 | 5551$\pm$21 | 0.55$\pm$0.04 |  $<$4 | 38
$\delta$ Pav | GJ 780 | 20:08:46.81 | -66:11:13.53 | G8IV | 6.09$\pm$0.01 | 0.99${}^{+0.14}_{-0.10}$ | 5590${}^{+156}_{-140}$ | 1.25${}^{+0.02}_{-0.04}$ | 6.2 | 1,2
LHS 2156 | HIP 47080 | 09:35:38.55 | +35:48:32.33 | G8IV/V | 11.20$\pm$0.02 | 0.96${}^{+0.16}_{-0.09}$ | 5499${}^{+156}_{-136}$ | 0.80${}^{+0.02}_{-0.04}$ |  $\cdots$ | 1,2
82 Eri | GJ 139, HD 20794 | 03:20:00.08 | -43:03:59.59 | G8V | 6.003$\pm$0.009 | 0.94$\pm$0.12 | 5413${}^{+91}_{-118}$ | 0.65$\pm$0.01 | 5.76; $>12$ | 1,2,9,13
$\tau$ Ceti | HD 10700, GJ 71 | 01:44:02.17 | -15:56:01.25 | G8.5V | 3.603$\pm$0.007 | 0.92${}^{+0.12}_{-0.10}$ | 5333${}^{+124}_{-75}$ | 0.51$\pm$0.01 |  $>$12 | 1,2,9
44 Boo Ba |  | 15:03:46.74 | +47:39:15.89 | K0V | 12.8$\pm$0.2 | 0.98 | 5300 | 0.51 |  $\cdots$ | 46,47,48
55 Cnc A |  $\rho^{1}$ Cnc, HIP 43587 | 08:52:35.22 | +28:19:47.22 | K0IV/V | 12.59$\pm$0.01 | 0.90${}^{+0.14}_{-0.09}$ | 5252${}^{+123}_{-172}$ | 0.64$\pm$0.02 | 9.5 | 1,2
70 Oph A | LHS 458, HIP 88601A | 18:05:27.47 | +02:29:42.81 | K0V | 5.15$\pm$0.03 | 0.89$\pm$0.02 | 5300$\pm$50 | 0.65$\pm$0.02 | 6.2$\pm$1.0 | 1,22
40 Eri A | GJ 166A, HD 26965 | 04:15:13.91 | -07:40:05.08 | K0.5V | 5.04$\pm$0.01 | 0.85${}^{+0.11}_{-0.10}$ | 5092${}^{+140}_{-149}$ | 0.43$\pm$0.01 | 6.9$\pm$4.7 | 1,14,23
$\delta$ Eri | HIP 17378 | 03:43:14.80 | -09:45:36.30 | K1III/IV | 9.04$\pm$0.07 | 1.19 | 5095 | 3.37${}^{+0.07}_{-0.06}$ | 6.28 | 2
GL 451A | Groombridge 1830 | 11:53:04.16 | +37:41:34.18 | K1IV | 9.16${}^{+0.16}_{-0.17}$ | 0.66 | 4950 | 0.23$\pm$0.04 | 5.0$\pm$0.3 | 2,6
GL 117 | HIP 13402, HD 17925 | 02:52:32.56 | -12:46:14.00 | K1.5V | 10.355$\pm$0.005 | 0.89${}^{+0.09}_{-0.13}$ | 5225${}^{+91}_{-136}$ | 0.40$\pm$0.01 | 1.5;$<$1.2;0.1 | 1,2,6,9
$\alpha$ Cen B | GJ 59 B, LHS 51 | 14:39:35.06 | -60:50:15.10 | K2IV | 1.347$\pm$0.03 | 0.87$\pm$0.07 | 5234$\pm$63 | 0.52$\pm$0.05 | 8, $>$11 | 2,9,30
GL 783 A | HD 191408A, LHS 486 | 20:11:12.54 | -36:06:29.58 | K2.5V | 6.05$\pm$0.03 | 0.74 | 4922 | 0.28$\pm$0.10 | 7.05$\pm$0.65 | 1,2,6,25
GL 183 | HD 32147, HIP 23311 | 05:00:49.59 | -05:45:30.96 | K3V | 8.845$\pm$0.003 | 0.81${}^{+0.11}_{-0.09}$ | 4931${}^{+133}_{-134}$ | 0.29$\pm$0.01 | 2 | 1,3,6
GL 892 | HD 219134 | 23:13:21.06 | +57:10:10.80 | K3V | 6.531$\pm$0.004 | 0.79$\pm$0.03 | 4817$\pm$62 | 0.29$\pm$0.01 | 11.0$\pm$2.2,12.46 | 1,2,36
GJ 667 A | HIP 84709 | 17:18:56.53 | -34:59:25.40 | K3V | 6.97$\pm$0.83 | 0.59 | 4810 | 0.13${}^{+0.17}_{-0.09}$ |  $\cdots$ | 2,3,31
LHS 1875 | GL 250, HIP 32984 | 06:52:17.47 | -05:10:25.42 | K3.5V | 8.747$\pm$0.004 | 0.76${}^{+0.09}_{-0.10}$ | 4719${}^{+157}_{-129}$ | 0.22${}^{+0.02}_{-0.01}$ | 3.1,$<$0.48 | 1,2,9
GL 570 A | HIP 73184 | 14:57:29.18 | -21:25:23.31 | K4V | 5.881$\pm$0.003 | 0.75${}^{+0.08}_{-0.10}$ | 4681${}^{+165}_{-126}$ | 0.22${}^{+0.01}_{-0.02}$ |  $<$0.6;3 | 1,2,9
$\xi$ Boo B |  | 14:51:23.19 | +19:06:03.97 | K4V | 6.71$\pm$0.16 | 0.66$\pm$0.07 | 4350$\pm$150 | 0.10$\pm$0.02 |  $<$4 | 38
44 Boo Bb |  | 15:03:46.74 | +47:39:15.89 | K4V | 12.8$\pm$0.2 | 0.55 | 5035 | 0.24 |  $\cdots$ | 46,47,48
GJ 667 B | WDS J17190-3459 B |  $\cdots$ |  $\cdots$ | K4V |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | 31
70 Oph B | LHS 459 | 18:05:27.82 | +02:29:39.12 | K5V | 5.15$\pm$0.03 | 0.73$\pm$0.01 | 4390$\pm$200 | 0.15$\pm$0.02 | 6.2$\pm$1.0 | 1,22
61 Cygni A | HIP 104214 | 21:06:59.64 | +38:45:49.90 | K5V | 3.497$\pm$0.001 | 0.67${}^{+0.08}_{-0.09}$ | 4304${}^{+161}_{-128}$ | 0.15$\pm$0.01 | 6 | 1,2,26
61 Cygni B | HIP 104217 | 21:07:00.88 | +38:45:21.86 | K7V | 3.497$\pm$0.001 | 0.61${}^{+0.08}_{-0.09}$ | 3949${}^{+156}_{-137}$ | 0.104${}^{+0.012}_{-0.009}$ | 6 | 1,2,26
GL 412A | HIP 54211 | 11:05:22.09 | +43:31:51.41 | M1V | 4.83$\pm$0.03 | 0.38 | 3688 | 0.020${}^{+0.090}_{-0.014}$ | 3 | 2,15,28
GL 570 B | (HIP 73182) | 14:57:27.70 | -21:25:07.94 | M1.5V | 5.86$\pm$0.06 | 0.55$\pm$0.05 | 3345 | 0.13$\pm$0.09 |  $\cdots$ | 1,2,37
GJ 832 | LHS 3685, HD 204961 | 21:33:33.90 | -49:00:45.47 | M1.5V | 4.964$\pm$0.001 | 0.44$\pm$0.02 | 3601 | 0.03${}^{+0.09}_{-0.02}$ | 9.24 | 1,2,4
GJ 667 C |  | 17:19:00.30 | -34:59:52.06 | M1.5V | 7.244$\pm$0.005 | 0.33$\pm$0.02 | 3350$\pm$50 | 0.0137$\pm$0.0009 |  $>$2 | 31
Kapteyn’s Star | GJ 191 | 05:11:50.38 | -45:02:37.73 | M2V | 3.9331$\pm$0.0004 | 0.28$\pm$0.02 | 3742$\pm$157 | 0.016$\pm$0.004 | 11.5${}^{+0.5}_{-1.5}$ | 1,2,16
Wolf 1055 | HD 180617, GL 752 A | 19:16:55.26 | +05:10:08.04 | M2.5V | 5.912$\pm$0.02 | 0.48$\pm$0.02 | 3534$\pm$51 | 0.032$\pm$0.001 |  $\cdots$ | 40
AD Leo | GJ 388 | 10:19:35.70 | 19:52:11.30 | M3V | 4.964$\pm$0.001 | 0.423$\pm$0.012 | 3477$\pm$23 | 0.0236$\pm$0.0001 | 0.4$\pm$0.1 | 39
GL 570 C | HIP 73182 | 14:57:27.70 | -21:25:07.94 | M3V | 5.86$\pm$0.06 | 0.39$\pm$0.03 |  $\cdots$ | 0.018 |  $\cdots$ | 37
Wolf 1061 | GJ 628, LHS 419 | 16:30:17.96 | -12:40:04.26 | M3.5V | 4.306$\pm$0.001 | 0.30$\pm$0.02 | 3309$\pm$157 | 0.011${}^{+0.003}_{-0.002}$ |  $\cdots$ | 1,2
Luyten’s Star | GJ 273 | 07:27:25.11 | +05:12:33.78 | M3.5V | 5.92$\pm$0.02 | 0.29$\pm$0.02 | 3382$\pm$49 | 0.0087${}^{+0.0063}_{-0.0065}$ |  $\cdots$ | 1,2,43
55 Cnc B |  | 08:52:40.28 | +28:18:54.91 | M4.5V | 12.48$\pm$0.02 | 0.26$\pm$0.02 | 3187$\pm$157 | 0.008$\pm$0.002 |  $\cdots$ | 1,35
GL 783B | HD 191408B, LHS 487 | 20:11:12.80 | -36:06:32.17 | M4.5V | 6.05$\pm$0.03 | 0.24 |  $\cdots$ | 0.0009 | 7.05$\pm$0.65 | 25
GJ 777B | LHS 3509 | 20:03:37.57 | +29:53:45.49 | M4.5V | 15.97$\pm$0.01 | 0.21$\pm$0.02 | 3099$\pm$157 | 0.0047$\pm$0.0012 |  $\cdots$ | 1,34
40 Eri C | LHS 25 | 04:15:19.12 | -07:40:15.25 | M4.5V | 5.04$\pm$0.01 | 0.204$\pm$0.006 | 3100 | 0.008 | 6.9$\pm$4.7 | 1,14,23,24
$\upsilon$ And B |  | 01:36:50.16 | 41:23:25.98 | M4.5V | 13.47$\pm$0.02 | 0.19$\pm$0.02 | 3159$\pm$157 | 0.0044${}^{+0.0012}_{-0.0011}$ |  $\cdots$ | 29
Proxima Centauri | GJ 551 | 14:29:34.16 | -62:40:33.87 | M5.5V | 1.3012$\pm$0.0003 | 0.13$\pm$0.02 | 2900$\pm$100 | 0.0016${}^{+0.0004}_{-0.0006}$ | 4.85 | 1,17,18
GL 412B | WX UMa, LHS 39 | 11:05:24.50 | +43:31:33.27 | M6.6V | 4.83$\pm$0.03 | 0.095 | 2700 | 0.00095 | 3 | 27,28
VB 10 | GL 752 B | 19:16:57.61 | +05:09:01.59 | M8V | 5.912$\pm$0.02 | 0.0881${}^{+0.0026}_{-0.0024}$ | 2508${}^{+63}_{-60}$ | 0.000499±0.000004 | 1 | 41,42

References. — 1: Stassun et al. (2019), 2: ExoCat-1, Turnbull (2015), 3: Gray
et al. (2006), 4: Sanz-Forcada et al. (2010), 5: Huang et al. (2018), 6:
Mamajek & Hillenbrand (2008), 7: Benedict et al. (2022), 8: Gaia Collaboration
(2020), 9: Takeda et al. (2007), 10: Gray et al. (2003), 11: Delrez et al.
(2021), 12: Udry et al. (2019), 13: Pepe et al. (2011), 14: Ma et al. (2018),
15: Mann et al. (2015), 16: Anglada-Escude et al. (2014), 17: Faria et al.
(2022), 18: Kervella et al. (2003), 19: Ghezzi et al. (2010), 20: Dorren &
Guinan (1994), 21: Güdel et al. (1997), 22: Eggenberger et al. (2008), 23:
Mason et al. (2017), 24: Johnson & Wright (1983), 25: Turnbull & Tarter
(2003), 26: Kervella et al. (2008), 27: Casagrande et al. (2008), 28: Mann et
al. (2015), 29: Rosenthal et al. (2021), 30: Santos et al. (2013), 31:
Anglada-Escudé et al. (2013), 32: Bidelman (1985), 33: Feng et al. (2019), 34:
Hawley et al. (1996), 35: Alonso-Floriano et al. (2015), 36: Gillon et al.
(2017), 37: Mariotti et al. (1990), 38: Fernandes et al. (1998), 39:
Kossakowski et al. (2022), 40: Burt et al. (2021), 41: Pravdo & Shaklan
(2009), 42: Pineda et al. (2021), 43: Astudillo-Defru et al. (2017), 44:
Holmberg et al. (2009), 45: Mosser et al. (2008), 46: Zirm (2011), 47:
Latković et al. (2021), 48: Zasche et al. (2009), 49: Philipot et al. (2023),
50: Nordström et al. (2004)

∗Significant disagreements in literature are discussed in Appendix A.

Table 3: Summary of X-ray Observations

| XMM-Newton |  | Chandra
---|---|---|---
Star | # Obs. | Exp. Time (ks) |  | # Obs. | Exp. Time (ks)
$\eta$ Crv | 0 |  $\cdots$ |  | 4 | 38.5
$\xi$ Oph | 0 |  $\cdots$ |  | 1 | 19.8
$\upsilon$ And | 1 | 5.3 |  | 4 | 58.9
$\iota$ Hor | 32 | 211.3 |  | 0 |  $\cdots$
$\nu$ Phe | 1 | 16.0 |  | 0 |  $\cdots$
$\gamma$ Pav | 1 | 15.9 |  | 0 |  $\cdots$
$\beta$ Vir | 1 | 41.1 |  | 0 |  $\cdots$
GJ 1095 | 0 |  $\cdots$ |  | 1 | 96.2
LHS 237 | 1 | 15.6 |  | 0 |  $\cdots$
$\iota$ Per | 0 |  $\cdots$ |  | 1 | 4.9
⋮ | ⋮ | ⋮ |  | ⋮ | ⋮
Total | 151 | 2235.7 |  | 41 | 815.0

Note. — Select entries are shown to illustrate the table form and content. The
full machine-readable table (57 rows) is available online from the journal.

###  3.2 Chandra

We retrieved Chandra ACIS imaging observations from the Chandra
archive777https://cda.harvard.edu/chaser/ and reprocessed all evt1 data using
the CIAO v4.15 (Fruscione et al., 2006) task chandra_repro and standard
reduction procedures. We ran the point source detection algorithm wavdetect on
each individual exposure to generate a list of X-ray source positions and
error ellipses. The major and minor axes of these error ellipses were
increased by a factor of five, and all X-ray sources were masked so that a
background light curve could be extracted for each observation. We inspected
these background light curves for background flaring events. We created good
time intervals (GTIs) using lc_clean and filtered all event data on these
GTIs. Finally, we restricted the energy range of our resulting “clean” evt2
files to 0.5-7 keV.

|
---|---
Figure 2: Comparison of the spatial resolution of XMM-Newton (left) and
Chandra (right); a 1′ scale bar is shown in both images. Left: The GL 412
binary star system with the locations of GL 412A (M1V) and GL 412B/WX Ma
(M6.6V) shown by cyan circles. The circles have radii of 10′′ and are
separated by $\sim$30′′. A third, unidentified X-ray source (dashed yellow
circle) is located $\sim$37′′ from GL 412B. Right: The 40 Eri triple star
system with the locations of 40 Eri A (K0V), 40 Eri B (DA white dwarf), and 40
Eri C (M4.5V) circled in white. The separation between the B and C components
is $\sim$7′′. X-ray sources are found coincident with 40 Eri A and C, while 40
Eri B is undetected.

We re-ran wavdetect on the cleaned image to identify significant X-ray sources
and obtain first-pass estimates of their source properties. Due to its
exquisite spatial resolution (0.496′′ pixel size) and a typical on-axis PSF
size less than 1′′ for the ACIS instruments, Chandra is the instrument of
choice for resolving crowded fields. For stellar systems with multiple
components and angular separations of less than a few arcseconds, Chandra is
currently the only telescope that can be used to definitively associate X-ray
emission with a given stellar component. Figure 2 (right panel) illustrates
the spatial resolving power of Chandra with a 5 ks exposure of the triple star
system 40 Eri, where the brightest X-ray source is clearly associated with the
position of 40 Eri C, a fainter X-ray source is found at the location of 40
Eri A, and 40 Eri B is undetected.

If the target star was detected in the cleaned evt2 image with $\gtrsim$50 net
counts, we extracted source light curves (binned to 100 s) to search for
evidence of X-ray variability. For sources with $\gtrsim$500 net counts, the
CIAO tool specextract was used to extract spectra and generate RMFs and ARFs.
All spectra were binned to contain at least 10 counts per bin.

###  3.3 Swift

The X-Ray Telescope (XRT) onboard the Neil Gehrels Swift Observatory
additionally observed some nearby stars. Swift is optimized to quickly slew to
observe transient events, and typically obtains short “snapshot” X-ray
observations ($\sim$2 ks or less) with XRT of relatively bright X-ray sources.
These observations are useful for providing flux estimates (or flux upper
limits) for stars of interest, but often cannot provide the longer-duration
light curves or higher signal-to-noise spectra that can be obtained from XMM-
Newton or Chandra. We therefore use the High Energy Astrophysics Science
Archive Research Center (HEASARC)888https://heasarc.gsfc.nasa.gov/cgi-
bin/W3Browse/swift.pl to determine which stars were observed by Swift but not
XMM-Newton or Chandra, and we use the Swift data products
generator999https://www.swift.ac.uk/user_objects/index.php to create stacked
images of these additional stars and measure count rates (or count rate upper
limits). The 19 stars that were observed only by Swift/XRT, but not by XMM-
Newton or Chandra, are not considered part of our primary stellar sample, but
we still provide updated provide X-ray luminosity estimates in Appendix B.

##  4 Analysis Procedure

We first inspected all X-ray observations for evidence of X-ray emission at
the location of the target stars. For stars that were detected, we first
extract light curves (binned to 100 s) to search for rapid variability within
the observation. For brighter stars (i.e., those with $\gtrsim$2000 net counts
with XMM-Newton or $\gtrsim$500 net counts with Chandra), changes in observed
count rate are readily apparent by visual inspection; for fainter stars, a
more quantitative assessment of variability is needed in addition to visual
inspection.

###  4.1 Assessing Count Rate Variability

We use the Anderson-Darling test, which is a non-parametric statistical test
that is more sensitive to short duration variations than the more commonly
used Kolmogorov-Smirnov (KS) test (Feigelson et al., 2022). The Anderson-
Darling statistic, $A^{2}$, is computed as (Rahman & Chakrobartty, 2004):

$A^{2}=-N-\frac{1}{N}\sum_{i=1}^{N}[(2i-1)\text{ln}(t_{i})+(2N+1-2i)\text{ln}(1-t_{i})],$
(1)

where $N$ is the number of photons observed and $t_{i}$ is the arrival time of
the $i$-th bin, scaled to the interval (0,1). We use scipy’s anderson routine
to compute $A^{2}$ and the critical values for the 5% and 1% significance
levels when compared to a normal distribution. An $A^{2}$ value that is larger
than a critical value for a given significance level is evidence of
variability; the typical 5% critical value is $\sim$0.75 and the typical 1%
critical value is $\sim$1.05 for the observations considered here. We
additionally compare the observed light curve to a constant count rate and
compute a standard reduced $\chi^{2}_{r}$ ($\chi^{2}$ divided by the degrees
of freedom) statistic.

Variability in observed count rate is, to first order, driven by a change in
X-ray flux incident on the detector. In our case, this change in flux is due
to changes in the intrinsic X-ray luminosity of the target star. However,
changes in the underlying spectra of the star can additionally impart more
subtle changes on the light curve, as the sensitivity of XMM-Newton and
Chandra to X-ray photons is energy-dependent. We therefore “flag” segments of
each light curve according to the type of variability observed, hereafter
referred to as the VarFlag, so that the spectral properties can be measured
and an accurate count rate-to-luminosity conversion factor can be established
as a function of time during the observation. The VarFlags are defined as
follows:

  * •

Q: quiescent periods, when the count rate is low and constant within the
uncertainties

  * •

F: strong flaring events, where the count rate is above $\sim$80% the
quiescent maximum count rate

  * •

D: decaying periods following flares or descending count rates that follow
highly elevated count rates; these periods are roughly defined as immediately
following an obvious flare where the count rate is below 80% the maximum flare
count rate but still above $\sim$150% of the quiescent count rate

  * •

E: elevated count rate periods, when count rates are approximately 50% higher
than the quiescent count rate but not obviously associated with a strong
flaring event

  * •

R: rising count rate periods that precede a strong flare

Dividing a single observation into a series of sub-exposures is a subjective
process: we wish to define sub-exposures that are short enough to capture
unique variations seen in the light curve (such as rapid flaring events, some
lasting only 200-300 s) but long enough so that a sufficient number of X-ray
counts are available for adequate spectral modeling (the $\sim$2000 and
$\sim$500 count limits for XMM-Newton and Chandra, respectively). Our main
goal is to accurately measure the X-ray luminosity of the star as a function
of time.

Figure 3 shows an example light curve of Proxima Centauri (XMM-Newton
observation number 0801880501) and a cumulative distribution function (CDF) of
photon arrival times compared to a constant count rate. The $A^{2}$ statistic,
critical values, and $\chi^{2}_{r}$ are shown, and the light curve has been
color-coded according to the VarFlag assigned to each segment of the light
curve. Table 4 summarizes the variability metrics for every X-ray observation
of every star available in the archives (the full machine-readable table is
available online from the journal). While some stars in our sample have been
the targets of extensive observing campaigns by XMM-Newton and/or Chandra for
many years, observational data is quite sparse for other stars. For stars with
many observations, long-term periodic X-ray variability has been detected and
is likely the result of coronal activity cycles (e.g., $\iota$ Hor, the
$\alpha$ Cen system; Sanz-Forcada et al., 2019; Ayres, 2023). Similar long-
term X-ray variability cycles may exist for sparsely observed stars in our
sample, but will not be detectable in single snapshot observations. Additional
observations are required to constrain the long-term coronal activity cycle
and/or flaring activity for many stars in the present sample.

\begin{overpic}[width=433.62pt]{ProxCen/ProxCen_0801880501_countrate_lc.png}
\put(59.0,8.0){\large F} \put(67.0,8.0){\large D} \put(80.0,8.0){\large Q}
\put(93.0,8.0){\large F} \end{overpic} Figure 3: Left: Photon arrival time CDF
of Proxima Centauri (XMM-Newton observation ID 0801880501, black) compared to
a constant count rate (dark red, dashed). The $A^{2}$ statistic and critical
values are shown in the upper left corner. Right: The light curve data (black
circles) compared to a constant count rate (dark red, dashed), with sections
of the light curve color-coded by VarFlag. Quiescent periods are shown in dark
blue-violet (“Q”), flaring periods in pink “F”), and descending count rates in
purple (“D”). Table 4: X-ray Variability Metrics

|  | Exp. | Net |  |
---|---|---|---|---|---
Star | Observation ID | Time (ks) | Counts | $\chi^{2}_{r}$ | $A^{2}$
$\iota$ Hor | XMM/0673610201 | 5.3 | 4230$\pm$70 | 1.5 | 0.49
$\beta$ Vir | XMM/0044740201 | 41.1 | 18450$\pm$160 | 3.1 | 0.73
$\upsilon$ And A | XMM/0722030101 | 5.3 | 940$\pm$40 | 2.1 | 0.36
44 Boo | XMM/0100650101 | 18.8 | 287110$\pm$550 | 3.1 | 0.64
GL 620.1A | XMM/0822070201 | 7.8 | 12220$\pm$110 | 1.9 | 0.41
$\kappa^{1}$ Cet | Chandra/22327 | 14.0 | 6600$\pm$80 | 1.7 | 3.26
Wolf 1061 | Chandra/20163 | 38.3 | 250$\pm$40 | 0.9 | 29.78
Prox Cen | XMM/0801880501 | 20.0 | 66790$\pm$260 | 35.2 | 17.55
GL 412B | XMM/0742230101 | 10.3 | 32030$\pm$210 | 243.1 | 3.43
VB 10 | XMM/0504010101 | 24.2 | 660$\pm$110 | 1.1 | 11.01
⋮ | ⋮ | ⋮ | ⋮ | ⋮ | ⋮

Note. — Select entries are shown to illustrate the table form and content. The
full machine-readable table (200 rows) is available online from the journal.
Net counts are reported for 0.2-15 keV for XMM-Newton observations and 0.5-7
keV for Chandra observations.

###  4.2 Spectral Modeling

All X-ray spectral modeling was performed with XSPEC v12.11.1 (Arnaud, 1996)
using standard $\chi^{2}$ statistics. Depending on the number of counts and
the overall complexity of the spectra, we used one-, two-, or three-
temperature apec models to parameterize the X-ray spectra. Since all of the
stars are located very close to the Sun (within $\sim$20 pc), the intervening
hydrogen column density is expected to be low; we adopt a fixed absorbing
column $N_{\rm H}$ (using the tbabs model) of $10^{19}$ cm-2 for all stars (a
similar assumption was made by Brown et al., 2023). We assume abundances from
Wilms et al. (2000) for both the plasma emission and interstellar absorption
models. The plasma temperature and normalizations were left as free parameters
during the fit. The normalization of each apec component is defined as
(neglecting cosmological terms):

$\mathcal{N}=\frac{10^{-14}}{4\pi d^{2}}\int n_{e}n_{\rm H}dV,$ (2)

where $d$ is the distance to the source, $n_{e}$ and $n_{\rm H}$ are the
electron and hydrogen column densities (in units of cm-3), respectively, and
$dV$ is the volume element (in units of cm3). The integral $\int n_{e}n_{\rm
H}dV$ is the volume emission measure (VEM). Once the best-fit model is
obtained, we use the cflux convolution model to measure the flux on each
plasma component.

While the spectra of some stars required only one or two apec components to
achieve a statistically acceptable fit, most spectra required a third apec
component. We initialized all spectral models with the default (solar)
abundance. For some stars (e.g., $\iota$ Hor, $\kappa^{1}$ Cet) the best fit
model was insensitive to the exact choice of abundance; for these stars the
abundance was kept fixed at the solar value. If a statistically acceptable fit
was not obtained with solar abundances, we allowed the abundance to be a free
parameter of the fit, subject to the constraint that the abundance be the same
for all plasma components.

For stars that exhibited count rate variability, spectra were extracted and
modeled for each sub-exposure. Table 5 provides full best-fit parameters for
all stars with enough counts to perform spectral modeling for all observations
and sub-exposures (the full machine-readable table is available online from
the journal). For stars with multiple observations or multiple sub-exposures
with the same VarFlag, we present the average best-fit model for a given
VarFlag in Table 6. For uniformity, we refer to the coolest component of a
three-component model as “APEC #1,” the hottest component as “APEC #3,” and
the intermediate-temperature component as “APEC #2.” The two-temperature
thermal plasma models have their individual components matched to the three-
temperature component with the most similar temperature. Figure 4 shows an
example of the averaged spectrum for $\iota$ Hor, which did not exhibit
significant count rate variability in any of the publicly available 38 XMM-
Newton observations and was therefore assigned a VarFlag of Q (quiescent). The
best-fit models for the individual exposures are shown in black, with the
average best-fit spectrum superimposed in red.

Figure 4: Best-fit spectra for all 38 XMM-Newton observations of $\iota$ Hor
(black) with the average best-fit spectrum superimposed (red). Table 5: All
Best-Fit Parameters from X-ray Spectral Modeling

|  |  | APEC #1 |  | APEC #2 |  | APEC #3 |  |  |
---|---|---|---|---|---|---|---|---|---|---
Star | ObsID | VarFlag | $kT_{1}$ | $\mathcal{N}_{1}$ |  | $kT_{2}$ | $\mathcal{N}_{2}$ |  | $kT_{3}$ | $\mathcal{N}_{3}$ | log$L_{\rm X}$ | abund | $\left<\chi^{2}/{\rm dof}\right>$
|  |  | (keV) | ($10^{-4}$) |  | (keV) | (10-4) |  | (keV) | (10-4) | [erg s-1] |  |
$\eta$ Crv | 14474 | Q | 0.16${}^{+0.07}_{-0.14}$ | 3.3${}^{+10.0}_{-1.6}$ |  | 0.50$\pm$0.08 | 4.6${}^{+5.0}_{-2.9}$ |  |  $\cdots$ |  $\cdots$ | 28.43$\pm$0.06 | 0.32${}^{+0.56}_{-0.16}$ | 1.14
$\beta$ Vir | 0044740201 | Q | 0.09$\pm$0.01 | 4.7${}^{+1.2}_{-0.9}$ |  | 0.37$\pm$0.01 | 2.4$\pm$0.1 |  |  $\cdots$ |  $\cdots$ | 28.08$\pm$0.02 | 1 (fixed) | 2.48
$\chi^{1}$ Ori | 0111500101/0-10ks | E | 0.10$\pm$0.02 | 39.7${}^{+55.7}_{-19.9}$ |  | 0.49$\pm$0.03 | 52.9${}^{+10.5}_{-9.4}$ |  | 0.84${}^{+0.10}_{-0.09}$ | 12.1${}^{+7.8}_{-4.8}$ | 28.88${}^{+0.03}_{-0.04}$ | 0.31${}^{+0.07}_{-0.05}$ | 1.31
$\beta$ Com | 0148680101/0-5ks | Q | 0.07$\pm$0.02 | 23.5${}^{+114.8}_{-15.1}$ |  | 0.39$\pm$0.02 | 3.4$\pm$0.3 |  |  $\cdots$ |  $\cdots$ | 28.11$\pm$0.02 | 1 (fixed) | 1.55
44 Boo | 0100650101/0-6ks | D | 0.25$\pm$0.02 | 23.8${}^{+4.4}_{-6.6}$ |  | 0.50${}^{+0.13}_{-0.12}$ | 9.8${}^{+5.4}_{-3.6}$ |  | 0.91${}^{+0.03}_{-0.02}$ | 275.3${}^{+13.8}_{-15.7}$ | 28.67$\pm$0.01 | 0.09$\pm$0.01 | 1.45
GL 620.1A | 0822070301 | Q | 0.11${}^{+0.02}_{-0.01}$ | 9.5${}^{+5.3}_{-3.6}$ |  | 0.46$\pm$0.03 | 1.4${}^{+2.5}_{-2.3}$ |  | 0.77${}^{+0.05}_{-0.04}$ | 6.3${}^{+1.9}_{-1.7}$ | 28.82$\pm$0.03 | 0.43${}^{+0.09}_{-0.06}$ | 1.60
⋮ | ⋮ | ⋮ | ⋮ | ⋮ |  | ⋮ | ⋮ |  | ⋮ | ⋮ | ⋮ | ⋮ | ⋮

Note. — Select entries are shown to illustrate the table form and content. The
full machine-readable table (360 rows) is available online from the journal.

Figure 5 shows the VEM for each star (calculated from the best-fit
normalization, defined in Equation 2 above) as a function of temperature for
the coolest APEC component (circles), the intermediate APEC component
(squares), and the hottest APEC component (diamonds) using the measurements
summarized in Table 6. The points are color-coded by stellar spectral type: F
and G stars are shown in yellow, K stars in orange, and M stars in red. In
general, the temperatures and VEMs of M stars exhibit the largest spread in
coronal temperatures relative to more massive stars. The coronal temperatures
of FG and K stars are similar for the coolest and hottest APEC components,
although VEMs of FG stars are roughly an order of magnitude higher than those
of K stars at similar temperatures. The VEMs do not significantly change
within a given spectral type group, but the intermediate-temperature APEC
component of FG stars is on average hotter (by $\sim$2 MK) than the
intermediate APEC component of K stars.

The three-component thermal plasma model that generally describes the X-ray
emission from F and G stars is described by plasma temperatures of $\sim$1.3
MK, 5.1 MK, and 9.4 MK, with corresponding (log) VEMs of 55.2 [cm3], 55.4
[cm3], and 55.0 [cm3], respectively. For K stars, the temperatures become 1.2
MK, 3.3 MK, 10 MK, with (log) VEM values of 54.1 [cm3], 53.9 [cm3], and 53.6
[cm3], respectively. For these hotter stars, periods of elevated count rates
or flaring events different from the quiescent periods in that VEMs increased,
while the plasma temperatures remained mostly constant. This result is in
sharp contrast to the model fits for M stars, which showed more significant
variations in both temperature and VEM during flaring events.

Table 6: Average Best-Fit Parameters from X-ray Spectral Modeling

|  | APEC #1 |  | APEC #2 |  | APEC #3 |  |  |
---|---|---|---|---|---|---|---|---|---
Star | VarFlag | $kT_{1}$ | $\mathcal{N}_{1}$ |  | $kT_{2}$ | $\mathcal{N}_{2}$ |  | $kT_{3}$ | $\mathcal{N}_{3}$ | log$L_{\rm X}$ | abund | $\left<\chi^{2}/{\rm dof}\right>$
|  | (keV) | ($10^{-4}$) |  | (keV) | (10-4) |  | (keV) | (10-4) | [erg s-1] |  |
$\eta$ Crv | Q | 0.20${}^{+0.12}_{-0.09}$ | 2.3${}^{+4.3}_{-1.1}$ |  | 0.55${}^{+0.09}_{-0.07}$ | 3.1${}^{+2.2}_{-1.7}$ |  |  $\cdots$ |  $\cdots$ | 28.41${}^{+0.05}_{-0.10}$ | 0.27${}^{+0.61}_{-0.18}$ | 1.31
$\upsilon$ And A | Q | 0.29$\pm$0.03 | 0.8$\pm$0.1 |  |  $\cdots$ |  $\cdots$ |  |  $\cdots$ |  $\cdots$ | 27.57$\pm$0.03 | 1 (fixed) | 0.91
$\iota$ Hor | Q | 0.09$\pm$0.03 | 7.7$\pm$1.1 |  | 0.42$\pm$0.08 | 2.9$\pm$0.2 |  | 0.71$\pm$0.14 | 1.6$\pm$0.2 | 28.75${}^{+0.10}_{-0.14}$ | 1 (fixed) | 1.27
$\beta$ Vir | Q | 0.09$\pm$0.01 | 4.7${}^{+1.2}_{-0.9}$ |  | 0.37$\pm$0.01 | 2.4$\pm$0.1 |  |  $\cdots$ |  $\cdots$ | 28.08$\pm$0.02 | 1 (fixed) | 2.48
$\chi^{1}$ Ori | Q | 0.08$\pm$0.03 | 258$\pm$228 |  | 0.42$\pm$0.07 | 28.6$\pm$9.2 |  | 0.69$\pm$0.06 | 21.3$\pm$8.2 | 28.94${}^{+0.07}_{-0.10}$ | 0.20$\pm$0.12 | 1.55
⋮ | ⋮ | ⋮ | ⋮ |  | ⋮ | ⋮ |  | ⋮ | ⋮ | ⋮ | ⋮ | ⋮

Note. — Select entries are shown to illustrate the table form and content. The
full machine-readable table (76 rows) is available online from the journal.

Figure 5: Best fit VEMs and temperatures for thermal plasma components from
XSPEC modeling. VEMs and temperatures from APEC #1 are shown in circles, those
from APEC #2 are shown in squares, and those from APEC # are shown in
diamonds. Values measured for F and G stars are shown in yellow, K stars in
orange, and M stars in red. M stars exhibit the largest variations in both VEM
and temperature across all three APEC components.

###  4.3 X-ray Luminosity

In addition to insights about the temperatures and VEMs of coronal structures
associated with our target stars, the best-fit X-ray spectral models provide
us with a count rate-to-luminosity conversion factor that mitigates against
instrumental effects from the extracted light curves and gives us a clearer
view of the luminosity evolution of the target star. Furthermore, we can
“unfold” our X-ray spectra (i.e., de-convolve the observed spectrum and the
energy-dependent detector response function) to assess the intrinsic energy
emission from the target star as a function of energy (or wavelength).

Figure 6 shows the continuum luminosity of each star (in units of $L_{\odot}$)
as a function of wavelength during quiescent times. All spectra have been
corrected for minor foreground absorption, and we use the XSPEC command fakeit
to extend our best-fit model to a wavelength range of $\sim$1-200 Å. For
nearly all stars, flaring events are driven by (sometimes dramatic)
enhancements in X-ray luminosity at wavelengths $\lesssim$10 Å (energies
$\gtrsim$1.2 keV). Figure 7 shows quiescent and flaring spectra for three
representative stars ($\kappa^{1}$ Cet, GJ 570A, and Prox Cen) from 1-1000 Å
(top panel), as well as the ratio of the flaring and quiescent spectra (bottom
panel). For both $\kappa^{1}$ Cet (G5V) and GJ 570A (K4V), the flux at $\sim$1
Å ($\sim$12 keV) during flaring times is a factor of $\sim$30 larger than
during quiescent times, but the fluxes at longer wavelengths ($\gtrsim$100 Å;
energies $\lesssim$0.12 keV) are similar. In the case of Prox Cen (M5.5V),
however, the flaring spectrum is brighter than the quiescent spectrum by at
least an order of magnitude at all wavelengths, with short-wavelength
($\lesssim$ 3 Å) emission enhanced by more than a factor of 103. This finding
is unlikely to be the result of the reduced soft energy responses of Chandra
or XMM-Newton, as the loss of effective area due to contamination build-up
occurs most significantly at energies below 1 keV (e.g., Grant et al., 2024).
Nevertheless, we use detector response functions specific to the orbital cycle
in which the individual observations were obtained to de-convolve the detector
response as accurately as possible. Plots of individual stellar X-ray spectra,
color-coded by VarFlag, are discussed in Appendix A.

---
Figure 6: Average continuum quiescent spectra for each target star (0.5-150 Å,
corresponding to an energy range of $\sim$25-0.08 keV). Spectra have been
corrected for minor foreground absorption. Stars are color-coded by spectral
type, with earlier-type stars shown in yellow and later-type stars shown in
dark red. Figure 7: Top: Quiescent (solid lines) and flaring (dashed lines)
continuum spectra of three representative stars from our sample ($\kappa^{1}$
Cet in yellow, GJ 570 A in orange, and Prox Cen in red). Bottom: The ratio of
flaring spectrum to quiescent spectrum, showing a pronounced enhancement in
X-ray emission at $<$10 Å during flaring periods for all three stars. The
flaring spectrum of Prox Cen remains an order of magnitude brighter than the
quiescent spectrum even at longer wavelengths.

A question that our analysis can elucidate is: For what fraction of time is a
star observed to be emitting at a given X-ray luminosity? While this is
related to the flaring frequency and the typical peak luminosity and duration
of flares, our goal in this work is not to directly characterize flaring
events or the physical mechanisms that cause them. Rather, we wish to predict
the X-ray environments in which HZ planets reside.

To do this, we use our best fit spectral models to first convert the observed
count rates to luminosities. For stars that are too faint to have their
spectra directly modeled, we use the spectral model of the star with the
closest spectral type and physical parameters to convert the observed count
rates into luminosities. For stars that are detected but too faint for
reliable light curve extraction, we simply estimate their time-averaged X-ray
luminosity; for un-detected stars we calculate 3$\sigma$ luminosity upper
limits. We then apply Astropy’s (Astropy Collaboration et al., 2013, 2018)
bayesian_block algorithm to the light curves; the “fitness” parameter is used
to define the measurement uncertainties. The Bayesian block algorithm
(Scargle, 1998; Scargle et al., 2013) is a segmentation technique, which aims
to optimally split (in this case) time series data into “blocks” such that
each block is statistically different from its neighboring block. An example
Bayesian blocked light curve for $\kappa^{1}$ Cet is shown in Figure 8.

Figure 8: The light curve of $\kappa^{1}$ Cet from XMM-Newton observation
0111410101. The MOS1 count rates were converted to luminosities (gray circles)
using the best-fit spectral models from three sub-exposures (from 0-14 ks,
14-25 ks, and 25-39 ks). The Bayesian blocked light curve is superimposed in
blue.

The output of the bayesian_block routine is a list of time bins and the
luminosity the star was emitting at during each bin. From these lists, we
measure the fraction of the time that each star was observed to be fainter
than a given $L_{\rm X}$ or $L_{\rm X}$/$L_{\rm bol}$ value. The results are
shown in Figure 9. For comparison, $L_{\rm X}$ of the modern-day Sun is
6.31$\times 10^{26}$ erg s-1 (log$L_{\rm X}$/$L_{\rm bol}$ = -6.78) when quiet
and 8.04$\times 10^{27}$ erg s-1 (log$L_{\rm X}$/$L_{\rm bol}$ = -5.68) when
active (Linsky et al., 2020), while $\sim$4 Gyr ago the $L_{\rm bol}$ of the
Sun was $\sim$75% the modern value and its $L_{\rm X}$ was a factor of
$\sim$2-3 higher (Obridko et al., 2020). On this plot, a perfectly non-
variable star would be a straight vertical line, while stars that experience
significant changes in X-ray luminosity (i.e., Proxima Centauri) will show a
pronounced horizontal extent. We measure the luminosities at which the stars
spend less than 25%, 50%, 70%, and 90% of the observed time (which we refer to
as $L_{25}$, $L_{50}$, $L_{75}$, and $L_{90}$, respectively). These values are
summarized in Table 7; we also include the time-averaged luminosities (with
uncertainties) of faintly detected stars and the 3$\sigma$ upper limits for
non-detected stars in the $L_{50}$ column.

Figure 9: Top: The fraction of time that a star is observed to be fainter than
a specific luminosity. Bottom: The fraction of time that a star is observed to
be below a $L_{\rm X}$/$L_{\rm bol}$ value. Stars are color-coded by spectral
type, with earlier-type stars shown in yellow and later-type stars shown in
dark red. The 25th, 50th, 75th, and 90th percentile luminosities and $L_{\rm
X}$/$L_{\rm bol}$ ratios are indicated by blue dashed lines (values for each
star summarized in Table 7). The shaded gray region indicates typical ranges
for the current Sun, and the blue shaded regions indicate approximate ranges
for the early Sun ($\sim$4 Gyr ago; see text).

Table 7: X-ray Luminosities and $L_{\rm X}$/$L_{\rm bol}$ Ratios (or 3$\sigma$
Upper Limits) for All Sample Stars

|  log$L_{25}$ |  log$L_{50}$ |  log$L_{75}$ |  log$L_{90}$ |  log($L_{25}$/$L_{\rm bol}$) |  log($L_{50}$/$L_{\rm bol}$) |  log($L_{75}$/$L_{\rm bol}$) |  log($L_{90}$/$L_{\rm bol}$)
---|---|---|---|---|---|---|---|---
Star |  [erg s-1] |  [erg s-1] |  [erg s-1] |  [erg s-1] |  |  |  |
$\eta$ Crv | 28.19 | 28.39 | 28.47 | 28.59 | -6.06 | -5.86 | -5.78 | -5.66
$\xi$ Oph | 27.47 | 27.64 | 27.77 | 28.03 | -6.73 | -6.56 | -6.43 | -6.17
$\upsilon$ And A | 27.27 | 27.55 | 27.72 | 27.87 | -6.84 | -6.56 | -6.39 | -6.24
$\beta$ Vir | 27.90 | 27.99 | 28.05 | 28.15 | -6.26 | -6.17 | -6.11 | -6.01
$\beta$ Hyi | 26.23 | 26.53 | 26.94 | 27.37 | -7.92 | -7.62 | -7.21 | -6.78
$\kappa^{1}$ Cet | 28.63 | 28.71 | 28.78 | 28.88 | -4.89 | -4.81 | -4.74 | -4.67
HD 140901 | 27.62 | 27.92 | 28.05 | 28.23 | -6.05 | -5.75 | -5.62 | -5.44
$\delta$ Eri | 26.39 | 26.53 | 26.69 | 26.89 | -7.72 | -7.58 | -7.42 | -7.22
GL 570 A | 27.15 | 27.36 | 27.52 | 27.68 | -5.78 | -5.57 | -5.41 | -5.25
61 Cyg A | 26.75 | 26.92 | 27.08 | 27.28 | -6.00 | -5.84 | -5.68 | -5.56
61 Cyg B | 26.69 | 26.87 | 27.04 | 27.42 | -5.89 | -5.71 | -5.54 | -5.35
AD Leo | 28.42 | 28.51 | 28.69 | 29.36 | -3.54 | -3.45 | -3.27 | -2.98
Proxima Centauri | 26.70 | 26.93 | 27.71 | 28.63 | -4.09 | -3.86 | -3.08 | -2.33
⋮ | ⋮ | ⋮ | ⋮ | ⋮ | ⋮ | ⋮ | ⋮ | ⋮

##  5 Discussion

We compared the 25th and 90th percentile $L_{\rm X}$/$L_{\rm bol}$ ratios for
stars in our sample to the current minimum and maximum ratios for the Sun
(Linsky et al., 2020), and find that eight stars that were detected in the
X-ray observations at high significance exhibit $L_{\rm X}$/$L_{\rm bol}$
ratios similar to that of the modern Sun: $\xi$ Oph (F2V), $\beta$ Vir (F9V),
47 UMa (G0V), $\beta$ Com (G0V), 55 Cnc A (G8V), GL 892 (K3V), GL 783 (K3V),
and GL 183 (K3V). Three of these stars ($\beta$ Vir, $\beta$ Com, and GL 183)
are younger than the Sun, while four are older (47 UMa, 55 Cnc A, GL 892, and
GL 783) and one ($\xi$ Oph) does not have an age estimate available in the
literature. Twelve stars have $L_{\rm X}$/$L_{\rm bol}$ ratios consistent with
those of the early Sun: $\eta$ Crv (F2V), $\beta$ Com, GL 311 (G1.5V),
$\zeta^{1}$ Ret (G2V), GL 327 (G3V), HD 140901 (G7IV), the unresolved binary
70 Oph (K0V+K5V), LHS 1875 (K3V), GL 570A (K4V), 61 Cyg A (K5V) and B (K7V),
and the unresolved GL 570 BC pair (M1.5V+M3V). We note that one star, $\beta$
Com, exhibits a range in $L_{\rm X}$/$L_{\rm bol}$ that fits within the narrow
overlap region between the modern and early Sun (log$L_{\rm X}$/$L_{\rm
bol}$$\approx$ -6.4 to -5.7). An additional 11 stars were not detected or only
marginally detected and have $L_{\rm X}$/$L_{\rm bol}$ upper limits lower than
the modern day minimum solar $L_{\rm X}$/$L_{\rm bol}$: $\nu$ Phe, LHS 237,
LHS 208, $\iota$ Per, $\rho$ CrB, GL 672, 51 Peg, GJ 777A, 82 Eri, GL 451A,
and Kapteyn’s Star. For roughly a dozen stars in our sample, the observed
$L_{\rm X}$ (or $L_{\rm X}$ upper limit) translates into X-ray surface fluxes
below the $\sim$104 erg s-1 cm-2 minimum observed in the quiet Sun and other
“coronal hole” stars (Schmitt, 1997). A closer look at these very X-ray faint
stars will be the subject of an upcoming study (Binder et al., in
preparation). In Table 8, we summarize these X-ray-detected Solar-like stars
(their spectral types and approximate ages are taken from Table 3.1) and the
relevant solar $L_{\rm X}$/$L_{\rm bol}$ comparison (modern or early Sun). We
additionally provide the “tier” each star was assigned in Mamajek &
Stapelfeldt (2024), indicating the priority which each star should be
considered for an HWO direct imaging survey (with A indicating the highest
priority group, and tiers B and C requiring further study).

Table 8: Potential HWO Target Stars with Solar-Like $L_{\rm X}$/$L_{\rm bol}$ Ratios Star Name | Spectral Type | Age (Gyr) | Sun Comparison | HWO Tier
---|---|---|---|---
$\eta$ Crv | F2V |  $<$2 | early | C
$\xi$ Oph | F2V | unknown | modern | C
$\nu$ Phe | F9V |  $\sim$3 |  $<$modern min | B
$\beta$ Vir | F9V |  $\sim$3 | modern | C
LHS 237 | F9V | unknown |  $<$modern min | none
$\iota$ Per | F9.5V |  $\sim$4 |  $<$modern min | A
LHS 208 | G0V |  $\sim$3 |  $<$modern min | B
$\rho$ CrB | G0V |  $\sim$11 |  $<$modern min | B
GL 672 | G0V |  $\sim$12 |  $<$modern min | C
47 UMa | G0V |  $\sim$6.5 | modern | A
$\beta$ Com | G0V |  $<$2 | modern,early | A
GL 311 | G1.5V |  $<$2 | early | B
$\zeta^{1}$ Ret | G2V |  $<$2 | early | A
51 Peg | G2V |  $\sim$7 |  $<$modern min | none
GL 327 | G3V |  $<$2 | early | C
GJ 777A | G7IV/V |  $\sim$13 |  $<$modern min | B
HD 140901 | G7IV/V |  $\sim$3 | early | C
82 Eri | G8V |  $\sim$6 |  $<$modern min | B
70 Oph AB | K0V+K5V |  $\sim$6 | early | B
55 Cnc A | K0IV/V |  $\sim$9.5 | modern | C
GL 451A | K1IV |  $\sim$5 |  $<$modern min | C
GL 783 | K2.5V |  $\sim$7 | modern | B
GL 892 | K3V |  $\sim$11 | modern | A
GL 183 | K3V |  $\sim$2 | modern | B
LHS 1875 | K3.5V |  $<$3 | early | C
GL 570A | K4V |  $<$3 | early | A
61 Cyg AB | K5V+K7V |  $\sim$6 | early | A
GL 570 BC | M1.5V+M3V |  $<$3 | early | none
Kapteyn’s Star | M2V |  $\sim$11 |  $<$modern min | none

Any HZ planets that may exist around the stars listed in Table 8 are currently
experiencing high energy radiation environments that are similar to or less
hostile than the modern Earth. Out of the 229 unique stars we identified as
potential targets for future direct imaging surveys, only $\sim$25% had ever
been imaged with either XMM-Newton or Chandra (this fraction increases to
$\sim$30% when Swift imaging is included), and out of the 57 stellar systems
with high quality X-ray observation we find that 28 systems ($\sim$50%) have
$L_{\rm X}$/$L_{\rm bol}$ ratios similar to or less than that of either the
modern or early Sun. These stars span the full range of spectral types and
ages considered here, and represent all three tiers of the Mamajek &
Stapelfeldt (2024) ranking system. It is therefore likely that there are an
additional $\sim$100 potential direct imaging target stars with HZs with
current X-ray radiation environments similar to or even less hostile than that
of the Sun. We note that threshold $L_{\rm X}$/$L_{\rm bol}$ratios below that
of the modern or early Sun could be a conservative condition for finding
habitable planets. Although, the past evolution of the star would also have to
be taken into account. Detailed coupled geological and atmospheric evolution
analyses would be most informative for predicting the targets with the best
chances of maintaining potentially habitable conditions (e.g., Krissansen-
Totton & Fortney, 2022). A systematic X-ray survey of these as-of-yet
unobserved nearby stars is needed to inform the selection of the best targets
for direct imaging surveys.

In general, M stars exhibit significantly more X-ray variability than earlier
type stars and spend a significant amount of time at $L_{\rm X}$/$L_{\rm bol}$
ratios dramatically larger than that of the Sun. We can supplement our sample
of Sun-like and low mass stars with X-ray observations of an 23 additional M-
and K-type stars from the MUSCLES/Mega-MUSCLES survey (Brown et al., 2023).
Brown et al. (2023) found a major decrease in the X-ray-to-bolometric
luminosity ratio ($L_{\rm X}$/$L_{\rm bol}$) with increasing effective
temperature. We compute $L_{\rm X}$/$L_{\rm bol}$ (or upper limits, in the
case of X-ray non-detections) for all stars in our sample using $L_{50}$ from
Table 7 and $L_{\rm bol}$ from Table 3.1. In Figure 10, we plot the
(logarithmic) $L_{\rm X}$/$L_{\rm bol}$ ratio as a function of $T_{\rm eff}$
(also from Table 3.1) for our full sample (colored points), with the
MUSCLES/Mega-MUSCLES sample shown in gray. Our sample extends this inverse
relationship between $L_{\rm X}$/$L_{\rm bol}$ and $T_{\rm eff}$ by $\sim$1000
K in $T_{\rm eff}$, with $L_{\rm X}$/$L_{\rm bol}$ ratios that are about an
order of magnitude lower than observed for K and M stars.

There is a small clustering of stars with $T_{\rm eff}\sim 5200-6200$ K with
significantly higher $L_{\rm X}$/$L_{\rm bol}$ ratios than the majority of
comparable stars, likely due to young stellar ages. Chromospheric activity in
solar-type stars is observed to decrease rapidly after $\sim$2-3 Gyr (Pace &
Pasquini, 2004; Zhang et al., 2019), after which activity levels remain mostly
constant. We indicate the approximate age of each star in Figure 10: young
stars ($\lesssim$2 Gyr) are shown in red, older stars ($\gtrsim$2 Gyr) are
shown in purple, and stars without age estimates or ambiguous ages are shown
in dark blue. Most of these high-$L_{\rm X}$/$L_{\rm bol}$, high-$T_{\rm eff}$
stars are indeed believed to be consistent with the younger ages where each
spectral type emits elevated high energy emission (Richey-Yowell et al.,
2023). There are three stars with ambiguous ages that lie within this region
of the diagram ($\chi^{1}$ Ori, 55 Cnc A, and GL 451A); the observed X-ray
emission of these stars suggests that younger age estimates may be preferred
for these stars.

Figure 10: The $L_{\rm X}$/$L_{\rm bol}$ ratio as a function of $T_{\rm eff}$.
Stars in our sample with firm X-ray detections are shown in circles, while
upper limits are shown for stars that were not X-ray detected (downward
pointing triangles). Stars that are believed to be young ($\lesssim$2 Gyr) are
shown in red, older stars ($\gtrsim$2 Gyr) are shown in purple, and stars
without age estimates or with ambiguous/discrepant age estimates are shown in
dark blue. The MUSCLES/Mega-MUSCLES sample (Brown et al., 2023) is shown in
gray. The yellow circles and shaded area indicate the quiescent and flaring
$L_{\rm X}$/$L_{\rm bol}$ ratios for the Sun (Linsky et al., 2020).

In Figure 11 we again show the $L_{\rm X}$/$L_{\rm bol}$ ratio as a function
of $T_{\rm eff}$ for stars in our sample, but we use $L_{25}$, $L_{50}$,
$L_{75}$ and $L_{90}$ to compute the $L_{\rm X}$/$L_{\rm bol}$ ratio. This
allows us to assess the degree to which intrinsic X-ray variability
contributes to the scatter in $L_{\rm X}$/$L_{\rm bol}$ observed in Figure 10.
The ratio of $L_{75}/L_{25}$ provides an estimate of how much a star varies
“normal” about the median $L_{\rm X}$. We find that the variability of the
youngest stars in our sample exhibit the smallest dynamic range, with
$L_{75}/L_{25}\sim 1.3$ for stars with $T_{\rm eff}\gtrsim$5000 K (for the
coolest young star in the plot, AD Leo, $L_{75}/L_{25}\sim 1.8$). The
$L_{75}/L_{25}$ ratio for older stars shows more dispersion, varying by up to
a factor of $\sim$3.5.

Figure 11: Top: the $L_{\rm X}$/$L_{\rm bol}$ ratio as a function of $T_{\rm
eff}$ assuming $L_{25}$ (black), $L_{50}$ (blue), $L_{75}$ (green), and
$L_{90}$ (dark yellow). The yellow circles and shaded area indicate the
quiescent and flaring $L_{\rm X}$/$L_{\rm bol}$ ratios for the Sun (Linsky et
al., 2020). Bottom: the ratio of $L_{75}/L_{25}$ as a function of $T_{\rm
eff}$, color-coded by age (as in Figure 10: young stars shown in red, old
stars shown in purple, and stars with ambiguous ages shown in dark blue).

##  6 Conclusions

We have analyzed $\sim$3 Msec of high quality XMM-Newton and Chandra
observations of 57 nearby stellar systems, as well as snapshot Swift images of
an additional 19 stellar systems, to provide a uniform catalog of the X-ray
properties of potential target stars for future direct imaging surveys. All
X-ray light curves and spectroscopic data is publicly available on Zenodo
(DOI: 10.5281/zenodo.11490574). The X-ray environments in the HZs around F, G,
and early K stars are more likely to be similar to those that the Earth has
evolved in, and we identify 29 stars with current $L_{\rm X}$/$L_{\rm bol}$
ratios similar to or less than that of the Sun. Our results are consistent
with Richey-Yowell et al. (2023), who found that the X-ray fluxes of late M
stars can be a factor of $\sim$3-15 times larger than for K stars. In addition
to high average $L_{\rm X}$/$L_{\rm bol}$ ratios, M stars exhibit more
frequent and dramatic flares than earlier-type stars. We note that, of the 229
candidate target stars for direct imaging, only $\sim$30% have been observed
by XMM-Newton, Chandra, or Swift.

This work is supported by NASA Exoplanets Research Program (XRP) award
#80NSSC23K0039 (PI Turnbull). This research has made use of the NASA Exoplanet
Archive, which is operated by the California Institute of Technology, under
contract with the National Aeronautics and Space Administration under the
Exoplanet Exploration Program. B.B. is grateful for the warm hospitality of
the University of California, Santa Barbara Kavli Institute for Theoretical
Physics, where the majority of this manuscript was written; this research was
supported in part by grant no. NSF PHY-2309135 to the Kavli Institute for
Theoretical Physics. S.P. acknowledges support from NASA under award number
80GSFC21M0002. E.S. and S.P. acknowledge support from the CHAMPs (Consortium
on Habitability and Atmospheres of M-dwarf Planets) team, supported by the
National Aeronautics and Space Administration (NASA) under grant nos.
80NSSC21K0905 and 80NSSC23K1399 issued through the Interdisciplinary Consortia
for Astrobiology Research (ICAR) program. This work has made use of data from
the European Space Agency (ESA) mission Gaia
(https://www.cosmos.esa.int/gaia), processed by the Gaia Data Processing and
Analysis Consortium (DPAC,
https://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has
been provided by national institutions, in particular the institutions
participating in the Gaia Multilateral Agreement. This work made use of data
supplied by the UK Swift Science Data Centre at the University of Leicester.
We would like to thank Eric Mamajek for useful discussions related to this
work, and the anonymous reviewer for helpful comments that improved our paper.

##  Appendix A Discussion of Individual Stars

In this appendix we provide additional context for each stellar system (e.g,
when there is a notable discrepancy in the published literature with regards
to a given physical parameter) and summarize the known or candidate exoplanets
associated with each star. We then present the results of our X-ray analysis
for each stellar system. Table 9 summarizes the observations of faint/non-
detected stars, including if the star was detected in a given observation or
not, the count rate (or upper limit on the count rate) and detection
significance, and the “match star” that was used to convert count rates to
luminosities.

Figure 12 shows X-ray images of resolved and unresolved binary star systems
included in our study. Individual images are discussed further for each star
below. Figure Set 13 shows the light curves extracted for all stars from all
observations in which they were detected, with sub-exposures colored by
VarFlag. Figure Set 14 shows the average best-fit spectra for all sufficiently
bright stars. For stars that exhibit flaring events, we also show the ratio of
the flaring-to-quiescent spectra as a function of wavelength.

|  |
---|---|---
(a) | (b) | (c)
|  |
(d) | (e) | (f)
|  |
(g) | (h) | (i)
Figure 12: Images of resolved and unresolved binary systems with XMM-Newton
and Chandra. (a) XMM-Newton/PN image of $\upsilon$ And. The locations of the
F8V star (A) and the M4.5V star (B) are shown by cyan circles (both with radii
10′′). The stars are separated by $\sim$0.9′. (b) XMM-Newton image of GL
620.1A and the white dwarf WD 1620-391. No X-rays are detected at the location
of WD 1620-391 in either observation. (c) The $\alpha$ Cen binary system as
seen by the PN camera on XMM-Newton. The positions of the G2V star (A) and K1V
star (B) are shown in blue circles (both with radii of 5′′ and separated by
$\sim$18′′). (d) The $\xi$ Boo binary system as seen by the MOS2 camera on
XMM-Newton. The locations of the G8V star (A) and K4V star (B) are shown in
cyan circles (both with radii of 5′′). (e) The 70 Oph binary star system as
seen by the PN camera on XMM-Newton. The positions of the K0V star (A) and the
K5V star (B) are shown in blue circles (both with radii of 5′′ and separated
by $\sim$3.7′′). (f) The GL 783 binary star system as seen by the PN camera on
XMM-Newton. The positions of the K3V star (A) and the M4.5V star (B) are shown
in cyan circles (both with radii of 5′′ and separated by $\sim$2.6′′. (g)
Chandra image of the GL 570 system. The primary GL 570A (K4V) is easily
resolved from the binary red dwarf pair GL 570BC (M1.5V and M3V,
respectively). Cyan circles have a radius of 10′′. The brown dwarf GL 570D is
not detected in X-rays. (h) The 61 Cyg binary system as seen by the PN camera
on XMM-Newton. The positions of the K5V star (A) and K7V star (B) are shown in
cyan circles (both with a radius of 10′′ and separated by $\sim$12′′). (i)
Chandra image of Wolf 1055 and VB 10 (both circles have radii of 10′′ and are
separated by $\sim$1.2′).
Some images have been smoothed for display purposes only; white contours indicate the intensity distribution of X-ray counts on the detector. Figure 13: Left: Photon arrival time CDF for $\alpha$ Cen (black) compared to a constant count rate (dark red, dashed). The $A^{2}$ statistic and critical values are shown in the upper-left corner. Right: The light curve data (black circles) compared to a constant count rate (dark red, dashed). Quiescent periods are shown in blue-violet, flaring periods in pink, elevated count rate periods in blue, and descending periods in purple. The complete figure set (193 images) is available in the online journal. Figure 14: Top: The best-fit continuum spectra for $\alpha$ Cen color-coded by VarFlag. Quiescent periods are shown in dark blue, elevated count rate periods in light blue, flaring periods in pink, periods of descending count rates in purple, and periods of rising count rates in red. Bottom: The ratio of the flaring-to-quiescent spectrum as a function of wavelength. The complete figure set (18 images) is available in the online journal. Table 9: Summary of Faint Star Observations Star | Observation ID | Detected? |  Count Rate (ct s-1) |  Significance ($\sigma$) | Match Star
---|---|---|---|---|---
$\xi$ Oph | Chandra/27852 | Yes |  (8.1$\pm$1.0)$\times 10^{-3}$ | 8.0 |  $\eta$ Crv
$\nu$ Phe | XMM/0206540101 | No |  $<$4.3$\times 10^{-3}$ |  $\cdots$ |  $\iota$ Hor
$\gamma$ Pav | XMM/0670380101 | Yes |  (1.1$\pm$0.4)$\times 10^{-2}$ | 2.8 |  $\iota$ Hor
GJ 1095 | Chandra/4199 | Yes |  (2.1$\pm$0.7)$\times 10^{-4}$ | 2.9 |  $\iota$ Hor
LHS 237 | XMM/0840210501 | No |  $<$3.7$\times 10^{-2}$ |  $\cdots$ |  $\iota$ Hor
LHS 208 | XMM/0865400201 | No |  $<$1.4$\times 10^{-2}$ |  $\cdots$ | GL 620.1A
LHS 208 | Chandra/22293 | No |  $<$5.7$\times 10^{-4}$ |  $\cdots$ | GL 620.1A
$\iota$ Per | Chandra/12338 | No |  $<$1.2$\times 10^{-2}$ |  $\cdots$ |  $\iota$ Hor
$\rho$ CrB | Chandra/12396 | No |  $<$3.8$\times 10^{-3}$ |  $\cdots$ |  $\iota$ Hor
GL 672 | Chandra/12397 | No |  $<$2.1$\times 10^{-3}$ |  $\cdots$ | GL 620.1A
47 UMa | XMM/0304203401 | Yes |  (3.9$\pm$0.3)$\times 10^{-2}$ | 12.1 | GL 620.1A
$\beta$ Hyi | Chandra/12337 | Yes |  (5.1$\pm$0.6)$\times 10^{-2}$ | 12.5 | GL 620.1A
$\beta$ Hyi | XMM/0006010401 | Yes |  (5.6$\pm$0.4)$\times 10^{-2}$ | 14.0 | GL 620.1A
18 Sco | XMM/0303660101 | Yes |  (2.2$\pm$0.3)$\times 10^{-2}$ | 7.3 | GL 311
18 Sco | Chandra/12393 | Yes |  (5.5$\pm$1.1)$\times 10^{-4}$ | 5.0 | GL 311
$\mu$ Ara | XMM/0551021001 | No |  $<$1.6$\times 10^{-2}$ |  $\cdots$ | GL 327
$\mu$ Ara | XMM/0551023101 | No |  $<$2.2$\times 10^{-2}$ |  $\cdots$ | GL 327
HD 136352 | XMM/0884680201 | Yes |  (4.1$\pm$1.5)$\times 10^{-3}$ | 2.7 |  $\kappa^{1}$ Cet
51 Peg | XMM/0551020901 | No |  $<$2.3$\times 10^{-3}$ |  $\cdots$ |  $\kappa^{1}$ Cet
51 Peg | Chandra/10825 | No |  $<$3.1$\times 10^{-3}$ |  $\cdots$ |  $\kappa^{1}$ Cet
GJ 777A | XMM/0304201101 | No |  $<$1.2$\times 10^{-2}$ |  $\cdots$ |  $\delta$ Pav
GJ 777A | XMM/0304202601 | No |  $<$7.8$\times 10^{-3}$ |  $\cdots$ |  $\delta$ Pav
HD 140901 | Chandra/13769 | Yes |  (8.9$\pm$0.9)$\times 10^{-3}$ | 11.0 |  $\delta$ Pav
82 Eri | XMM/0670380601 | No |  $<$1.2$\times 10^{-2}$ |  $\cdots$ | 55 Cnc A
82 Eri | Chandra/22292 | No |  $<$8.1$\times 10^{-4}$ |  $\cdots$ | 55 Cnc A
GL 451A | Chandra/9931 | No |  $<$4.9$\times 10^{-4}$ |  $\cdots$ | GL 183
$\tau$ Cet | XMM/0670380501 | Yes |  (1.2$\pm$0.2)$\times 10^{-2}$ | 7.0 | 55 Cnc A
$\tau$ Cet | Chandra/1886 | Yes |  (2.4$\pm$0.3)$\times 10^{-3}$ | 8.0 | 55 Cnc A
$\delta$ Eri | XMM/0205720101 | Yes |  (7.8$\pm$1.1)$\times 10^{-3}$ | 7.3 | 55 Cnc A
40 Eri A | Chandra/13644 | Yes |  (1.4$\pm$0.2)$\times 10^{-2}$ | 7.7 | GL 117
GL 412 A | XMM/0742230101 | Yes |  (5.7$\pm$0.2)$\times 10^{-2}$ | 28.5 | GJ 832
GJ 667 C | Chandra/17317 | Yes |  (5.7$\pm$0.9)$\times 10^{-3}$ | 6.3 | GJ 832
GJ 667 C | Chandra/17318 | Yes |  (3.1$\pm$0.5)$\times 10^{-4}$ | 6.2 | GJ 832
Kapteyn | Chandra/merged | No |  $<$1.2$\times 10^{-4}$ |  $\cdots$ | Wolf 1055
Luyten | Chandra/20164 | Yes |  (3.5$\pm$0.7)$\times 10^{-3}$ | 5.0 | AD Leo
$\upsilon$ And B | XMM/0722030101 | No |  $<$3.2$\times 10^{-2}$ |  $\cdots$ | 40 Eri C
VB 10 | XMM/0504010101 | Yes |  (2.7$\pm$0.5)$\times 10^{-2}$ | 5.4 | GL 412B
VB 10 | Chandra/616 | Yes |  (1.2$\pm$0.8)$\times 10^{-3}$ | 1.5 | GL 412B
VB 10 | Chandra/7428 | Yes |  (1.6$\pm$0.4)$\times 10^{-3}$ | 4.0 | GL 412B

###  A.1 $\eta$ Crv

There are no significant discrepancies in the stellar physical parameters of
$\eta$ Crv. Although there are no known or candidate exoplanets, the system is
known to host two debris disks: one at a distance of $\sim$160-180 AU, and a
hotter disk at $\sim$3-7 AU (Wyatt et al., 2005, 2007; Smith et al., 2008).

$\eta$ Crv was observed four times with Chandra. All observations were taken
with ACIS-S in 1/4 subarray mode for $\sim 10$ ks. In all observations, $\eta$
Crv was detected with $\sim$1000 net counts (0.5-7 keV), but no significant
variability was observed in the X-ray light curves (see Figure Set 13). The
spectrum of $\eta$ Crv is well-described by a two-component thermal plasma
model, with no significant changes in the temperatures or normalizations
across the different observations.

###  A.2 $\xi$ Oph

There are no significant discrepancies in the stellar physical parameters of
$\xi$ Oph, and the star does not host any no known or candidate exoplanets.
$\xi$ Oph was observed once with Chandra/ACIS-S for 19.8 ks and robustly
detected with $\sim$160 counts (8$\sigma$ significance). There are hints of
low-level variability in the X-ray emission of $\xi$ Oph, but the relatively
low number of net counts prohibits spectral modeling.

###  A.3 $\upsilon$ And

$\upsilon$ And A is an F8V dwarf in a presumed binary system with a proper
motion companion $\upsilon$ And B, an M4.5V dwarf (also referred to as
$\upsilon$ And D). The two stars have an angular separation of $\sim$0.9′
(corresponding to $\sim$750 AU). There are no significant discrepancies in the
physical parameters of either star. $\upsilon$ And A hosts three confirmed
exoplanets ($\upsilon$ And b, c, and d) which were discovered via RV
measurements. The masses of $\upsilon$ And b, c, and d are $M_{\rm p}\sin i$
$\sim$0.67 $M_{\rm Jup}$, $\sim$0.67 $M_{\rm Jup}$, $\sim$2 $M_{\rm Jup}$, and
$\sim$4 $M_{\rm Jup}$, respectively, with orbital periods (orbital distances)
of $\sim$4.6 days (0.06 AU), $\sim$241 days (0.82 AU), and $\sim$1282 days
(2.5 AU), respectively (all planetary data taken from Rosenthal et al., 2021).

The $\upsilon$ And system was observed four times with Chandra/ACIS-S in 1/8
subarray mode (for $\sim$15 ks per observation) and once with XMM-Newton for
5.3 ks. $\upsilon$ And A was detected in all five observations with sufficient
counts to extract light curves (see Figure Set 13) and with enough counts in
Chandra observations 10976 and 10977 to enable spectral modeling. There is no
evidence for strong X-ray variability in any observation, and both Chandra
spectra were consistent with a single-temperature thermal plasma model with a
temperature of $\sim$0.3 keV ($\sim$3.4 MK; see Table 6). $\upsilon$ And B was
out of the field of view of the Chandra observations. However, the star was
within the field of view of the XMM-Newton PN detector but fell just on the
gap between two CCDs and was not detected (see Figure 12).

###  A.4 $\iota$ Hor

$\iota$ Hor is characterized either as a late F-type (F9V; Gray et al., 2006)
or early G-type (G0V; Turnbull, 2015) star. There is additional disagreement
on the age of the system: Turnbull (2015) quotes an age of 2.72 Gyr, while
Sanz-Forcada et al. (2010) suggest a much younger age of $\sim$0.47 Gyr (see
also the detailed analysis of Sanz-Forcada et al., 2019, where an age of 600
Myr is adopted). We adopt a spectral type of F9V and an age of 0.47 Gyr in
this work. $\iota$ Hor hosts one confirmed exoplanet that was first detected
via radial velocity (RV) measurements by Kürster et al. (2000). It is a giant
planet with $M_{\rm p}\sin i$$\sim$2.3 $M_{\rm Jup}$ and an orbital period of
$\sim$300 days (corresponding to an orbital distance of $\sim$0.9 AU; Stassun
et al., 2017).

$\iota$ Hor is one of the most thoroughly observed stars in our sample, with
32 publicly available XMM-Newton observations taken between 2011 and 2018. A
detailed study of the coronal activity and structure of the star is presented
in Sanz-Forcada et al. (2019), who combined the XMM-Newton monitoring with
observations from TESS and the STIS instrument on the Hubble Space Telescope.
They corroborate a 1.6 year X-ray activity cycle for the star, although
$\iota$ Hor does not exhibit the same type of dramatic changes in luminosity
as seen in other stars in our sample (see Figure 9). Our X-ray analysis of
$\iota$ Hor is consistent with Sanz-Forcada et al. (2019), although we do not
undertake a more extensive study of the abundances of the star. Light curves
of $\iota$ Hor are shown in Figure Set 13, and the best-fit spectrum is
described in Table 6.

###  A.5 $\nu$ Phe

$\nu$ Phe is characterized either as a late F-type (F9V; Gray et al., 2006) or
early G-type (G0V; Turnbull, 2015) star. There are no confirmed or candidate
exoplanets in the system, and the star was undetected in a 16 ks XMM-Newton
observation.

###  A.6 $\gamma$ Pav

$\gamma$ Pav is a known F-type star, with the specific spectral type ranging
from F9V (Gray et al., 2006) to F7V (Turnbull, 2015). There are significant
discrepancies in the estimated age of $\gamma$ Pav: Holmberg et al. (2009)
quote an age of $\sim$1 Gyr based on photometric observations and comparison
to stellar evolution models, while and asteroseismic analysis by Mosser et al.
(2008) yields an age of 7.25 Gyr. There are no known or candidate exoplanets
orbiting $\gamma$ Pav. The star was only marginally detected (2.8$\sigma$) in
a 25.9 ks XMM-Newton observation. The low X-ray luminosity of $\gamma$ Pav is
consistent with the older age estimate of Mosser et al. (2008).

###  A.7 $\beta$ Vir

There are no significant discrepancies in the stellar physical parameters for
$\beta$ Vir, and the star does not host any currently known or candidate
exoplanets. $\beta$ Vir was observed once with XMM-Newton (observation ID
0044740201) and was detected with $\sim$18,450 net counts on the PN camera
(thick filter). The X-ray light curve does not show evidence of significant
variability (see Figure Set 13). The spectrum is well described by a two-
temperature thermal plasma model (see Table 6).

###  A.8 GJ 1095

The spectral type of GJ 1095 is given as F9V by Gray et al. (2003) or G0V by
Turnbull (2015). There are no known or candidate exoplanets orbiting GJ 1095.
The star was only marginally detected (2.9$\sigma$) in a 96.2 ks Chandra
observation.

###  A.9 LHS 237

There are no significant discrepancies in the stellar physical parameters of
LHS 237, and no age estimate is currently available in the literature. LHS 237
is a triple-star system: the inner spectroscopic binary has an orbital period
of $\sim$10 years (the angular separation of $\sim$0.3′′ is unresolvable by
XMM-Newton; Tokovinin, 2012), and the third component is one of the coolest
known white dwarfs (NLTT 18141 = GJ 288B) located $\sim$14.5′ from the central
binary (Holberg et al., 2013). There are no known or candidate exoplanets in
the system.

LHS 237 was serendipitously imaged by XMM-Newton for 15.6 ks in observation
0840210501 (the intended target for the observation was a Wolf-Rayet star).
The inner spectroscopic binary is not detected, and the white dwarf is outside
the XMM-Newton field of view.

###  A.10 $\iota$ Per

There are no significant discrepancies in the stellar physical parameters for
$\iota$ Per. The spectral type is either F9.5V (Gray et al., 2003) or G0V
(Turnbull, 2015). There are no known or candidate exoplanets in the system.
$\iota$ Per was observed for 4.9 ks with Chandra in ACIS-S 1/4 subarray mode,
but the star was not detected.

###  A.11 $\beta$ Hyi

There are no significant disagreements in the stellar physical parameters of
$\beta$ Hyi, and the star does not host any currently known or candidate
exoplanets. $\beta$ Hyi was observed once with Chandra/ACIS-S in 1/4 subarray
mode and once with XMM-Newton. The star was detected in both observations with
sufficient counts to extract light curves, but the star is too faint for
accurate spectral modeling to be performed. There is no strong evidence for
X-ray variability in either the Chandra or XMM-Newton light curves of $\beta$
Hyi (see Figure Set 13).

###  A.12 $\beta$ Com

There are only minor differences in the reported age of $\beta$ Com: Takeda et
al. (2007) report an age $<$1.12 Gyr, while Turnbull (2015) report an age of
1.7 Gyr. All other spectral parameters are taken from Stassun et al. (2017).
There are currently no known or candidate exoplanets orbiting $\beta$ Com.

$\beta$ Com was robustly detected (with $\sim$38,000 net counts) in a 57.2 ks
observation by XMM-Newton. We extracted a light curve from this observation
and found marginal evidence for count rate variability (see Figure Set 13).
Given the large number of counts available, we split the observation into 5
ks-long sub-exposures and extracted spectra from each sub-exposure to search
for evidence of underlying spectroscopic variability. Throughout the
observation, a two-temperature thermal plasma model was sufficient to describe
the X-ray spectrum of $\beta$ Com, with the best fit temperatures and
normalizations remaining constant within the uncertainties. All best-fit sub-
exposure spectral parameters for $\beta$ Com are provided in Table 5, and the
average best-fit spectral parameters are listed in Table 6. The best-fit
spectral models are shown in Figure Set 14.

###  A.13 LHS 208 ($\pi$ Men)

All spectral parameters for LHS 208 are taken from Huang et al. (2018). LHS
208 hosts three confirmed exoplanets. HD 390931 b is a giant planet that was
detected via RV measurements (Jones et al., 2002); it has $M_{\rm p}\sin
i$$\sim$12 $M_{\rm Jup}$ and an orbital period of $\sim$2090 days
(corresponding to an orbital distance of $\sim$3.3 AU; Feng et al., 2022).
$\pi$ Men c is a super-Earth that was discovered via transits (Gandolfi et
al., 2018); it has an $M_{\rm p}\sin i$$\sim$3.5 $M_{\oplus}$ with an orbital
period of $\sim$6 days ($\sim$0.07 AU orbital distance; Feng et al., 2022).
$\pi$ Men d was detected via RV measurements with $M_{\rm p}\sin i$$\sim$13
$M_{\oplus}$ and an orbital period of $\sim$125 days (Hatzes et al., 2022).
LHS 208 not detected in a 36.9 ks observation with XMM-Newton or a 19.6 ks
observation with Chandra.

###  A.14 $\rho$ CrB

Some minor discrepancies, mass ranges from 0.95 (Brewer et al., 2023) to 1.05
$M_{\odot}$ (Stassun et al., 2017) and spectral type from G0V (Gray et al.,
2006) to G2V (Turnbull, 2015). $\rho$ CrB hosts three confirmed exoplanets
($\rho$ CrB c, d, and e), with the existence of a fourth exoplanet
controversial ($\rho$ CrB b, Brewer et al., 2023, which may simply be the
result of correlated stellar activity in the RV measurements;). The masses of
$\rho$ CrB c, d, and e are $M_{\rm p}\sin i$ $\sim$28 $M_{\oplus}$, $\sim$22
$M_{\oplus}$, and 3.8 $M_{\oplus}$, respectively, with orbital periods of
$\sim$102 days ($\sim$0.4 AU), $\sim$282 days ($\sim$0.8 AU), and $\sim$13
days ($\sim$0.11 AU), respectively. $\rho$ CrB is also approaching the end of
its main sequence lifetime, with a prognosis of planetary engulfment for most
of the known planets in the system (Kane, 2023). $\rho$ CrB was observed once
with Chandra/ACIS-S for 9.8 ks, but was not detected.

###  A.15 GL 672

There are no significant discrepancies in the stellar physical parameters for
GL 672, and there are no known or candidate exoplanets in the system. GL 672
was observed for 9.9 ks by Chandra/ACIS-S but was not detected.

###  A.16 $\chi^{1}$ Ori

$\chi^{1}$ Ori is a G0V-type star in a $\sim$14 year elliptical ($e\sim 0.45$)
orbit with a low-mass companion. RV measurements compiled by Han & Gatewood
(2002) were used to infer a secondary mass of $\sim$0.15 $M_{\odot}$ with a
likely spectral type of M4V or M5V. The estimated angular semimajor axis of
the orbit is $\sim$0.668′′, unresolvable by XMM-Newton but potentially
spatially resolvable by Chandra (see Section 3). The inferred spectral
parameters may be biased due to the presence of this low-mass companion. There
is significant disagreement in the literature about the age of the system.
Turnbull (2015) and Takeda et al. (2007) estimate an age of 4.32 Gyr, while
Mamajek & Hillenbrand (2008) use activity-rotation diagnostics to estimate an
age of 300-400 Myr. There are no confirmed or candidate exoplanets in the
system.

Since only one XMM-Newton observation is available for $\chi^{1}$ Ori, we are
unable to identify the source of the X-ray emission – whether it be from the
G0V star, the low-mass companion, or both. There is one potential, minor
flaring event during the observation, starting $\sim$21 ks after the beginning
of the observation and lasting for $\sim$4 ks, where the peak count rate
exceeds the median count rate by $\sim$40% (see Figure Set 13). While this may
not be a bona fide flare event, we classify this period as flaring so that its
spectrum can be compared to the remainder of the observation. The light curve
otherwise shows a count rate that is either constant within the uncertainties
(which we classify as quiescent) or somewhat elevated above the median count
rate level by $\sim$20% (which we classify as elevated).

Figure Set 14 shows the quiescent, elevated, and potential flare spectra for
the $\chi^{1}$ Ori system. The elevated and flare spectra are nearly
indistinguishable, and both show excess emission over the quiescent spectrum
at wavelengths $\lesssim$10Å. This difference in short-wavelength emission is
driven primarily by the hottest thermal plasma component in our model, which
has a temperature of $\sim$0.8 keV for the elevated and flaring spectra but
$\sim$0.7 keV in the quiescent spectrum (see also Table 6). The integrated
X-ray luminosity of $\chi^{1}$ Ori is constant within the uncertainties across
the three spectra.

###  A.17 GL 788

There are no significant discrepancies in reported stellar parameters of GL
788. The system does not host any known or candidate exoplanets. GL 788 was
detected in a 9.6 ks XMM-Newton observation with $\sim$4470 net counts. There
was no evidence of count rate variability apparent in the X-ray light curve
(see Figure Set 13). The spectrum was well-described with a single, cool
thermal plasma ($kT\sim$0.2 keV).

###  A.18 47 UMa

There are no significant discrepancies in the stellar physical parameters
reported in the literature for 47 UMa. The star hosts three confirmed giant
exoplanets that were detected via radial velocity measurements. 47 UMa b, c,
and d have masses of $M_{\rm p}\sin i$$\sim$2.4 $M_{\rm Jup}$, $\sim$0.5
$M_{\rm Jup}$, and 1.5 $\sim$2.4 $M_{\rm Jup}$, respectively, and have orbital
periods (distances) of $\sim$1077 days ($\sim$2 AU), $\sim$2287 days
($\sim$7.8 AU), and $\sim$19000 days ($\sim$12.4 AU), respectively. All
planetary parameters were taken from Rosenthal et al. (2021).

47 UMa was detected at $\sim$12$\sigma$ significance in a 6.2 ks XMM-Newton
observation. We extracted a light curve from the observation but found no
evidence of variability (see Figure Set 13), and the star was too faint to
enable spectroscopic modeling.

###  A.19 GL 620.1A

GL 620.1A is either co-moving with (Mugrauer, 2019) or in a wide binary with
(Holberg et al., 2002) the DA-type white dwarf WD 1620-391. The angular
distance between the two components is $\sim$5.8′, which would correspond to a
physical separation of $\sim$4500 AU at the distance of GL 620.1A. GL 620.1A
is classified as a barium star, and its over-abundance in $s$-process elements
has been attributed to pollution due mass transfer during the white dwarf
progenitor’s AGB phase (Porto de Mello & da Silva, 1997). There is some
disagreement in the literature over the spectral type and age of GL 620.1A.
Spectral types range from G1V (Gray et al., 2006) to a later G3/5V (Turnbull,
2015); we adopt a spectral type G1V. While Turnbull (2015) quote an age of 2
Gyr, other estimates of the age of GL 620.1A are lower: Ghezzi et al. (2010)
and Takeda et al. (2007) derive upper limits of $<$1 Gyr and $<$0.68 Gyr,
respectively, while Sanz-Forcada et al. (2010) adopt an age of 0.4 Gyr. One
giant exoplanet is confirmed via RV measurements of GL 620.1A (Mayor et al.,
2004); the planet has $M_{\rm p}\sin i$$\sim$1.2 $M_{\rm Jup}$ and an orbital
period of $\sim$528 days (orbital distance of $\sim$1.3 AU).

GL 620.1A is detected in both publicly available XMM-Newton observations
(Figure 12) with sufficient counts to enable spectral modeling (see Figure Set
14 and Table 6). There is no significant variability present in the count rate
light curves of GL 620.1A (Table 4, see Figure Set 13), so we do not attempt
to divide the observations into sub-exposures.

####  A.19.1 The White Dwarf WD 1620-391

While GL 620.1A is robustly detected in both publicly available XMM-Newton
observations, WD 1620-391 is not (see Figure 12a). The 3$\sigma$ upper limits
on the count rates of 0.014 ct s-1 (observation ID 0822070201) and 0.0236 ct
s-1 (observation ID 0822070301). We use WebPIMMs101010See
https://heasarc.gsfc.nasa.gov/cgi-bin/Tools/w3pimms/w3pimms.pl to convert the
observed XMM-Newton/PN count rates (both observations were taken with the
thick filter) to a flux assuming an APEC thermal plasma model with $kT\approx
0.5$ keV (the approximate best-fit model for the accreting white dwarf G
29-38, which is detected in X-rays by both XMM-Newton and Chandra; Cunningham
et al., 2022; Estrada-Dorado et al., 2023). We find flux (luminosity) upper
limits of 2.0$\times 10^{-14}$ erg s-1 cm-2 (4$\times 10^{26}$ erg s-1) and
3.4$\times 10^{-14}$ erg s-1 cm-2 (6.8$\times 10^{26}$ erg s-1) for
observations 0822070201 and 0822070301, respectively. These luminosity limits
are a factor of $\sim$5-7 higher than the observed luminosity of G 29-38;
targeted observations with a more sensitive X-ray instrument are needed to
determine if WD 1620-391 is emitting X-rays at a levels comparable to other
X-ray-detected white dwarfs.

###  A.20 GL 311 ($\pi^{1}$ UMa)

There are no significant discrepancies in the stellar physical parameters
reported in the literature for GL 311. The system does not host any currently
known or candidate exoplanets.

The variability metrics for GL 311 (see Table 4) suggest that count rate
variability is observed in the available XMM-Newton observation (observation
number 0111400101; see Figure Set 13). There are two minor flare-like events
in the light curve. The first begins at $\sim$21 ks after the beginning of the
observation, and reaches a peak count rate that is a factor of $\sim$2 higher
than the median count rate. This event lasts for approximately 8 ks and it
immediately followed by a second count rate increase, which exceeds the median
count rate by a factor of $\sim$1.6 and lasts for approximately 5 ks. The
post-flare count rate is elevated by $\sim$15% relative to the pre-flare count
rate. Spectroscopic modeling before, during, and after the flaring event show
that the temperatures of all three thermal components increased in temperature
during the flaring event, although a modest decrease in the normalization
(related to the VEM) during the flare led to a constant (within the
uncertainties) luminosity (Figure Set 14; see also Table 6).

Thus, although GL 311 clearly exhibits a degree of rapid X-ray variability (on
a timescale of a few hours), this activity does not have a significant affect
on the overall X-ray luminosity output of the star. The relative stability of
the star’s X-ray luminosity is also apparent in Figure 9, where the luminosity
curve for GL 311 is near-vertical.

###  A.21 $\alpha$ Cen

$\alpha$ Centauri A (G2V) and B (K1V) are members of the extensively studied
triple system that also includes Proxima Centauri (M5.5V; see Section A.56).
$\alpha$ Cen A and B are in a $\sim$80 year elliptical orbit (e$\sim$0.5;
Akeson et al., 2021). One candidate exoplanet, $\alpha$ Cen B b, was reported
(Dumusque et al., 2012) but was later found to be a false positive signal
(Rajpaul et al., 2016). There are currently no other known or candidate
exoplanets in the $\alpha$ Cen AB system.

The semi-major axis of the $\alpha$ Cen system is $\sim$18′′, but due to the
brightness of the stars there is considerable blending in the XMM-Newton
images (see Figure 12). We therefore do not attempt to extract X-ray products
(light curves, spectra) for each star individually. The $\alpha$ Cen system
exhibits significant X-ray variability (see Figure Set 13). The system is
well-described by a three-component thermal plasma model, and increases in the
observed count rates are driven largely by increases in the hottest plasma
temperature (from $\sim$0.6 keV during quiescent periods to $\sim$0.8-0.9 keV
during times associated with flaring events) and normalization (which
increases by a factor of $\sim$12). The spectra of $\alpha$ Cen are shown in
Figure Set 14 (see also Table 6). Due to the significant blending of the two
components, we do not use the quiescent $\alpha$ Cen spectrum as a baseline
for inferring other stellar luminosities.

###  A.22 44 Boo

44 Boo (also called $\iota$ Boo) is a triple star system with a G0V primary
(44 Boo A) in $\sim$210 year orbit with a spectroscopic binary composed of a
K0V (44 Boo Ba) dwarf and a K4V (44 Boo Bb) dwarf (Zasche et al., 2009; Zirm,
2011). 44 Boo A is separated from the spectroscopic binary by $\sim$0.2′′. The
44 Boo B binary has an orbital period of only $\sim$6.4 hours and is a
possible contact binary (Lu et al., 2001). There are no known or candidate
exoplanets in the 44 Boo system.

44 Boo is robustly detected with XMM-Newton in an 18.8 ks observation, with
count rate variations clearly seen in the light curve (see Figure Set 13). We
do not include the final $\sim$5 ks in our spectral analysis due to a
significant background flaring event during the observation. Our spectral
modeling revealed no significant differences between the minor flaring event
from $\sim$6-10 ks and the relatively constant count rate period from 15-20
ks. The spectrum of 44 Boo did change during periods of declining count rate
(the first 6 ks and from 10-15 ks) as the result changes in the hottest
thermal plasma component: $kT_{3}$ decreased from $\sim$2.1-2.2 keV during
quiescent/flaring times to $\sim$1.5 keV during periods of secular decline,
while the normalization of this component (related to the VEM) more than
doubled. These changes to APEC #3 result in a decrease in X-ray flux at the
shortest wavelengths ($\lesssim$10 Å); see Table 6 and Figure Set 14.

###  A.23 18 Sco

There are minor differences in the reported spectral type of 18 Sco: Turnbull
(2015) quote a spectral type of G5V, but Gray et al. (2006) classifies the
star as a G2V. We adopt a spectral type of G2V. 18 Sco hosts one candidate
super-Earth with $M_{\rm p}\sin i$ $\sim$6.8 $M_{\oplus}$ and an orbital
period of $\sim$20 days (Laliotis et al., 2023). The star is detected in both
a 20.1 ks XMM-Newton observation and a 62.2 ks observation with
Chandra/ACIS-S. However, the 18 Sco is not bright enough in either observation
for timing or spectroscopic analysis to be performed.

###  A.24 51 Peg

There are no significant discrepancies reported for the stellar parameters of
51 Peg. The star is famous for hosting the first discovered exoplanet (Mayor &
Queloz, 1995), the hot Jupiter 51 Peg b with $M_{\rm p}\sin i$ $\sim$0.46
$M_{\rm Jup}$ and an orbital period of $\sim$4.2 days ($\sim$0.05 AU Rosenthal
et al., 2021). 51 Peg was observed with both Chandra and XMM-Newton, but was
not detected in either observation.

###  A.25 $\zeta$ Ret

The wide binary $\zeta$ Ret contains $\zeta^{1}$ Ret, which is listed as a
G3V/G5V by Turnbull (2015) but G2V as Gray et al. (2006), and $\zeta^{2}$ Ret,
a G2V star (del Peloso et al., 2000). The two stars are separated by
$\sim$5.2′. The system does not hot any known or candidate exoplanets.

Although this separation would be easily resolvable by XMM-Newton, only
$\zeta^{1}$ Ret is contained with the fields of view of the XMM-Newton
cameras. The star is robustly detected (with $\sim$7700 net counts) in a 20.9
ks XMM-Newton observation (ID 0404920101). No significant variability is
observed in the X-ray light curve of $\zeta^{1}$ Ret (see Figure Set 13), and
the spectrum is well-described by a two-component thermal plasma model (with
temperatures of $\sim$0.1 keV and $\sim$0.4 keV; see Figure Set 14 and Table
6).

###  A.26 HD 136352

There are no significant discrepancies in the stellar physical parameters
reported in the literature for HD 136352. The star hosts three confirmed
exoplanets, all of which have masses in the super-Earth/sub-Neptune regime. HD
136352 b, c, and d have $M_{\rm p}\sin i$ of $\sim$4.7 $M_{\oplus}$, $\sim$11
$M_{\oplus}$, and $\sim$8.8 $M_{\oplus}$, respectively, with orbital periods
(distances) of $\sim$11.6 days ($\sim$0.1 AU), $\sim$28 days ($\sim$0.17 AU),
and $\sim$107 days ($\sim$0.43 AU), respectively. The inner two planets were
found to transit the host star by Kane et al. (2020b), providing bulk
densities for the planets. Subsequent observations by Delrez et al. (2021)
revealed that the third, outer planet also transits the host star. HD 136352
is marginally detected ($\sim$2.7$\sigma$) in a 16.9 ks XMM-Newton observation
with a count rate of 0.0041$\pm$0.0015 ct s-1 with the PN camera and medium
filter (Table 9).

###  A.27 GL 327

There are no significant discrepancies in reported stellar parameters of GL
327 in the literature. The system does not host any currently known or
candidate exoplanets. GL 327 was robustly detected with $\sim$10,500 net
counts in a 23.5 ks exposure with XMM-Newton (observaton ID 0404920201). Count
rate variability was observed in the X-ray spectrum as shown in Figure Set 13,
so we split the observation into three sub-exposures and modeled the three
spectra independently. In all three sub-exposures, a single-temperature
thermal plasma component was sufficient to describe the observations. During
the quiescent period from $\sim$6-21 ks, the plasma temperature was $\sim$0.4
keV, which increased to $\sim$0.5 keV during the descending and flaring
periods. The normalization during periods of increased X-ray count rate
additionally increased by $\sim$20%. The spectra of GL 327 are shown in Figure
Set 14 (see also Table 6.

###  A.28 $\mu$ Ara

No significant discrepancies in stellar parameters are present in the
literature for $\mu$ Ara, although there is some uncertainty about the
evolutionary stage; we adopt a spectral type of G3IV/V (Turnbull, 2015). $\mu$
Ara hosts four known (giant) exoplanets, all detected via RV measurements. HD
160691 b, c, d, and e have masses of $M_{\rm p}\sin i$$\sim$1.7 $M_{\rm Jup}$,
$\sim$2.4 $M_{\rm Jup}$, $\sim$10.2 $M_{\oplus}$, and $\sim$0.5 $M_{\rm Jup}$,
respectively, and the planets have orbital periods (orbital distances) of
$\sim$646 days ($\sim$1.5 AU), $\sim$4470 days ($\sim$5.5 AU), $\sim$10 days
($\sim$0.09 AU), and $\sim$307 days ($\sim$0.9 AU), respectively. All
planetary data were taken from Goździewski et al. (2007). $\mu$ Ara is not
detected in either of the two publicly available XMM-Newton observations.

###  A.29 $\kappa^{1}$ Cet

There is some discrepancy in the reported age estimates of $\kappa^{1}$ Cet:
Turnbull (2015) lists an age of 2.2 Gyr and Takeda et al. (2007) derive an
upper limit of $<$2.76 Gyr, but many studies report significantly younger
ages. For example, Mamajek & Hillenbrand (2008) derive an age of 300-400 Myr,
while Güdel et al. (1997) estimated an age of 750 Myr and and Dorren & Guinan
(1994) estimate an age of 650 Myr. There are no currently known or candidate
exoplanets orbiting $\kappa^{1}$ Cet.

$\kappa^{1}$ Cet is relatively well-observed in X-rays, with five XMM-Newton
observations and two Chandra observations available in the archives, and the
star is easily detected in all exposures. It was also observed twice with
Chandra using ACIS-S7 in 1/8 subarray mode with a 0.44 s readout time in both
observations. While there are clear indications of variability in the light
curves of $\kappa^{1}$ Cet the maximum observed count rate only ever exceeds
the median count rate by a factor of $\sim$2 (see Figure Set 13). However,
despite the only modest changes in count rate, the spectra during these minor
flare-like events is quite distinct from the lower count rate quiescent
periods, with significant excess emission a wavelengths shortward of 10 Å
compared to quiescent times. The change in count rate is being driven by a
significant increase in the temperature (by $\sim$0.2 keV, or $\sim$2.3 MK)
and normalization (a factor of $\sim$2.2 increase) of the hottest component in
a three-thermal plasma component model (Figure Set 14; see also Table 6).

###  A.30 HD 140901

There are no significant discrepancies in the published stellar physical
parameters of HD 140901. Two planets have been identified via RV measurements.
HD 140901 c has $M_{\rm p}\sin i$ $\sim$1.8 $M_{\rm Jup}$ with an orbital
period of $\sim$14390 days (orbital distance of $\sim$11.8 AU; Philipot et
al., 2023). The existence of HD 140901 b is controversial; if the planet
exists, it has a mass of $M_{\rm p}\sin i$ of $\sim$16$M_{\oplus}$ with an
orbital period of $\sim$9 days (orbital distance 0.085 AU), but the RV data of
HD 140901 can be equally well described by a retrograde orbital inclination of
$\sim$138∘ (Feng et al., 2022; Philipot et al., 2023).

HD 140901 was observed once by Chandra/ACIS-I for 24.7 ks. The star was
robustly detected with $\sim$220 net counts. The light curve shows clear
evidence for variability (see Figure Set 13), but the star is not bright
enough to enable spectral modeling.

###  A.31 GJ 777

GJ 777 is a binary star system composed of a G6IV primary (GJ 777A) and an
M4.5V secondary (LHS 2509), which are separated by $\sim$3000 AU ($\sim$3′).
GJ 777A hosts two currently known exoplanets, detected via RV measurements: HD
190360 b has $M_{\rm p}\sin i$ $\sim$1.8 $M_{\rm Jup}$ and an orbital period
of $\sim$2854 days ($\sim$3.9 AU; Feng et al., 2021), and HD 190360 c has
$M_{\rm p}\sin i$ $\sim$19 $M_{\oplus}$ and an orbital period of $\sim$17 days
($\sim$0.13 AU; Rosenthal et al., 2021). No X-ray emission is detected
coincident with the GJ 777 binary system in either of two snapshot XMM-Newton
observations.

###  A.32 $\xi$ Boo

$\xi$ Boo is a visual binary ($\xi$ Boo A has a spectral type G8V and $\xi$
Boo B has a spectral type K4V; Turnbull, 2015) with an angular separation of
$\sim$4.9′′ (orbital period $\sim$150 years; Wielen, 1962), the stars are
unresolvable by XMM-Newton (see Figure 12). No discrepancies in stellar
physical parameters in the literature. No known or candidate exoplanets. $\xi$
Boo A is a known BY Draconis variable (Finley et al., 2019).

The $\xi$ Boo binary is robustly detected by XMM-Newton, although the two
components cannot be resolved. The system exhibits dramatic X-ray variability,
as shown in the light curve in Figure Set 13. The X-ray spectra of $\xi$ Boo
are described by three-component thermal plasma models in all variability
states. The flaring spectrum requires hotter temperatures for all three plasma
components than the quiescent spectrum, with the normalization of the hottest
component increasing by a factor of $\sim$3.9 during flaring times compared to
quiescent times. Figure Set 14 shows the best fit spectral models for $\xi$
Boo (see also Table 6).

###  A.33 $\delta$ Pav

There are no significant discrepancies in the stellar physical parameters
reported in the literature for $\delta$ Pav. The star does not host any
currently known or candidate exoplanets.

$\delta$ Pav was detected in a 32.6 ks XMM-Newton observation (ID number
0780510401) $\sim$1000 net counts. The $A^{2}$ statistic provides evidence
that the star may be variable on short timescales, however there are not a
sufficient number of counts to split the observation into sub-exposures (see
Figure Set 13). We attempted to model the spectrum of $\delta$ Pav despite the
low number of counts; the spectrum required two thermal plasma components to
achieve a statistically acceptable fit (see Table 6). The best-fit spectrum in
shown in Figure Set 14.

###  A.34 LHS 2156

There are no significant discrepancies in the stellar physical parameters
reported in the literature for LHS 2156. There are no known or candidate
exoplanets in the system.

LHS 2156 was detected with $\sim$60,000 net counts in a 68.4 ks XMM-Newton
observation with the PN camera (thick filter). Figure Set 13 shows the X-ray
light curve, which exhibits count rate variability over the duration of the
observation. The spectra of LHS 2156 are well described by three-component
thermal plasma models. The best-fit temperatures and normalizations of two
cooler plasma components remain constant within the uncertainties in all sub-
exposures. The hottest thermal plasma component increases in temperature by
$\sim$0.1 keV during flaring periods compared to quiescent periods, an the
normalization increases by $\sim$25%. The best-fit spectra for LHS 2156 are
shown in Figure Set 14 (see also Table 6).

###  A.35 82 Eri

There is some discrepancy in the literature about the age of 82 Eri: Turnbull
(2015) and Takeda et al. (2007) quote ages $>$12 Gyr, but a much younger age
of 5.76$\pm$0.66 Gyr was found by (Pepe et al., 2011). 82 Eri hosts four
confirmed exoplanets, all of which were discovered via RV measurements (Pepe
et al., 2011). HD 20794 b, c, d, and e have $M_{\rm p}\sin i$ values of
$\sim$2.8 $M_{\oplus}$, $\sim$2.5 $M_{\oplus}$, $\sim$3.5 $M_{\oplus}$, and
$\sim$4.8 $M_{\oplus}$, respectively, with orbital periods (distances) of
$\sim$18 days ($\sim$0.13 AU), $\sim$43 days ($\sim$0.23 AU), $\sim$89 days
($\sim$0.26 AU), and $\sim$147 days ($\sim$0.5 AU), respectively. All
planetary parameters were taken from Feng et al. (2017a). 82 Eri has been
observed by both XMM-Newton and Chandra (see Table 3), but was not detected in
either observation.

###  A.36 $\tau$ Ceti

There are no significant discrepancies in the stellar physical parameters of
$\tau$ Ceti reported in the literature. $\tau$ Ceti hosts four exoplanets, all
confirmed via radial velocity measurements (Feng et al., 2017b), with some
evidence of additional planets in or near the $\tau$ Cet HZ (Dietrich & Apai,
2021). $\tau$ Ceti e and f both have masses $M_{\rm p}\sin i$ $\sim$3.9
$M_{\oplus}$ with orbital periods of $\sim$163 days ($\sim$0.55 AU) and
$\sim$636 days ($\sim$1.3 AU), respectively, while $\tau$ Ceti g and h both
have masses of $M_{\rm p}\sin i$ $\sim$1.8 $M_{\oplus}$ and orbital periods of
$\sim$20 days ($\sim$0.13 AU) and $\sim$49 days ($\sim$0.24 AU), respectively.

$\tau$ Ceti was observed by Chandra ACIS-S for 45.1 ks in 1/8 subarray mode
and detected at $\sim$7$\sigma$ significance with $\sim$110 net counts. We
extracted a light curve from this observation and find an $A^{2}$ statistic
that is strongly suggestive of variability (see Table 4 and Figure Set 13),
but due to the relatively low number of counts we did not attempt to extract a
spectrum. $\tau$ Ceti was marginally detected in an 11.9 ks XMM-Newton
observation, but with insufficient counts for light curve extraction or
spectral modeling to be performed. The luminosity estimates for $\tau$ Ceti
suggest that the star’s X-ray luminosity changes by a factor of $\sim$2.

###  A.37 55 Cnc

55 Cnc is a binary star system containing a G8V star (55 Cnc A) and a M4.5V
dwarf (55 Cnc B). No significant discrepancies in stellar physical parameters.
The two components are separated by $\sim$1,065 AU ($\sim$1.4′; Raghavan et
al., 2006). Five exoplanets are known to orbit 55 Cnc A. 55 Cnc b was the
first to be discovered (Butler et al., 1997) with a mass of $M_{\rm p}\sin i$
$\sim$0.8 $M_{\rm Jup}$ and an orbital period $\sim$14.7 days (corresponding
to an orbital distance of $\sim$0.11 AU; Bourrier et al., 2018; Nelson et al.,
2014). 55 Cnc c, d, e, and f has masses of $M_{\rm p}\sin i$ $\sim$0.16
$M_{\rm Jup}$, $\sim$3.9 $M_{\rm Jup}$, $\sim$8 $M_{\oplus}$, and $\sim$0.15
$M_{\rm Jup}$, respectively, with orbital periods of $\sim$44 days ($\sim$0.24
AU), $\sim$4870 days ($\sim$5.5 AU), $\sim$0.7 days ($\sim$0.02 AU), and
$\sim$260 days ($\sim$0.8 AU), respectively.

Cnc 55 was observed twice by Chandra/ACIS-S in 1/8-subarry mode (for 10.7 ks
and 18 ks) and once for 10.4 ks by XMM-Newton. 55 Cnc A was not detected by
Chandra, but was detected (at $\sim$5.6$\sigma$) significance by XMM-Newton.
No evidence for strong variability was observed in the XMM-Newton light curve
of 55 Cnc A (Table 4, see also Figure Set 13). Despite the relatively low
number of net counts, we attempted to model the X-ray spectrum of 55 Cnc A. We
find that a single-component thermal plasma model is sufficient to describe
the X-ray spectrum (see Table 6 and Figure Set 14). 55 Cnc B was not detected
in the XMM-Newton observation or one of the Chandra observations (the star was
not within the field of view of the second Chandra observation).

###  A.38 70 Oph

70 Oph is a visual and spectroscopic binary containing a K0V (70 Oph A) star
and a K5V (70 Oph B) star with an orbital period of $\sim$88 years (Pourbaix,
2000). The orbit is highly elliptical, with an eccentricity $e=0.50$ that
causes the orbital separation of the stars to vary from $\sim$11.4 AU to
$\sim$34.8 AU (Eggenberger et al., 2008). These orbital separations correspond
to angular separations of $\sim$2.3-6.8′′, unresolvable by XMM-Newton but
potentially spatially resolvable by Chandra (see Section 3). Eggenberger et
al. (2008) find a best-fit age of 6.2$\pm$1.0 Gyr, in broad agreement with the
6.8-7.0 Gyr found by Tang et al. (2008). These ages (which where derived using
astroseismic observations) are in contrast to the significantly younger age
($\sim$680 Myr) predicted by Mamajek & Hillenbrand (2008). There are no
currently known or candidate exoplanets in the system.

There are three XMM-Newton observations available for the 70 Oph system. The
individual stars are not resolvable in the XMM-Newton observations (see Figure
12); we therefore cannot separate the X-ray emission from the K0V star (70 Oph
A) from the K5V star (70 Oph B) with the currently available observations.

Despite the changes in X-ray count rate observed in the light curves (see
Figure Set 13), the best-fit spectral models of 70 Oph across different
VarFlags are nearly indistinguishable (Figure Set 14; see also Table 6), and
the X-ray luminosity of the two stars changes only by $\sim$40% between the
flaring and quiescent periods. Since we cannot separate the X-ray emission of
the binary components individually, we do not use the best-fit spectrum of the
70 Oph system as template for inferring luminosities of other stars in our
sample.

###  A.39 40 Eri

The 40 Eri triple system contains a K0V star (40 Eri A) that is either in a
wide orbit or co-moving with a white dwarf (40 Eri B)-M4.5V dwarf (40 Eri C)
pair. The projected angular distance between 40 Eri A and 40 Eri BC is
$\sim$1.3′. While the orbital period of the wider A-BC pair is very long, the
orbital period of the B-C pair is measured to be $\sim$230 years (Mason et
al., 2017; Bond et al., 2017). The angular separation between 40 Eri B and 40
Eri C is $\sim$6.9′′, easily resolvable by Chandra (see Figure 2). 40 Eri C
has a visual magnitude of $\sim$11, which we use to estimate $L_{\rm bol}$,
and Johnson & Wright (1983) estimate $T_{\rm eff}$ to be $\sim$3,100 K. There
is one confirmed planet orbiting 40 Eri A, which was discovered via RV
measurements. HD 26965 b has $M_{\rm p}\sin i$ of $\sim$8.5 $M_{\oplus}$ with
an orbital period of $\sim$42 days (Ma et al., 2018). No currently known or
candidate planets are associated with the 40 Eri BC pair.

The first soft (0.1-2.0 keV) X-ray detection of the 40 Eri system was obtained
by HEAO-1, which measured a luminosity of (9.6$\pm$3.2)$\times 10^{28}$ erg
s-1 (corrected for different distance assumed in that work; Cash et al.,
1979). Due to the poor spatial resolution of HEAO-1, the binary components
could not be resolved. Different scenarios for X-ray production are explored
in Cash et al. (1979), and the X-ray emission was attributed to 40 Eri C.

The Chandra image of the system clearly shows X-ray sources coincident with
both 40 Eri A and C (although not the white dwarf, 40 Eri B; see discussion
below). 40 Eri C is the brighter of the two sources, with roughly an order of
magnitude more net counts than observed for 40 Eri A. There is no evidence for
strong variability in either star (see Figure Set 13). While the coolest
thermal plasma component in the 40 Eri C spectrum is similar in temperature to
other late-type stars in our sample, the hotter two thermal plasmas are
considerably hotter (see Table 6 and Figure Set 14). Despite the lack of
obvious variability in 40 Eri C, the best-fit X-ray spectral parameters most
closely resemble those of WX UMa in its flaring and post-flaring states. It is
possible this star indeed exhibits significant coronal variability, but it was
not detectable in the short snapshot Chandra observation available in the
archive.

####  A.39.1 The White Dwarf 40 Eri B

The white dwarf 40 Eri B is not detected in the 5 ks Chandra/ACIS-S
observation of the system. The 3$\sigma$ upper limit on the count rate is
0.0033 ct s-1. We again use WebPIMMs to convert the observed count rate upper
limit to a flux (and luminosity) upper limit, assuming the same spectral model
as for WD 1620-391 (Section A.19.1). We find a flux upper limit of 1.6$\times
10^{-13}$ erg s-1 cm-2, which corresponds to a luminosity of $\sim$5$\times
10^{25}$ erg s-1. This luminosity upper limit is below the measured X-ray
luminosity of G 29-38 ($\sim$8$\times 10^{25}$ erg s-1; Cunningham et al.,
2022; Estrada-Dorado et al., 2023), strongly suggesting a lack of a cold
accretion disk around this white dwarf.

###  A.40 $\delta$ Eri

There are no significant discrepancies in the reported physical parameters for
$\delta$ Eri, and the star does not host any known or candidate exoplanets.
$\delta$ Eri was observed once with XMM-Newton (observation observation) for
56.2 ks. There was marginal evidence for significant variability in the X-ray
light curve of $\delta$ Eri (see Figure Set 13), but the star was too faint
for spectral modeling.

###  A.41 GL 451A (Groombridge 1830)

There are no significant discrepancies in the stellar physical parameters
reported in the literature for GL 451A. The system does not host any currently
known or candidate exoplanets. GL 451A was observed once with Chandra/ACIS-S
(observation ID 9931) for 32.8 ks in 1/8-subarray mode across ACIS detectors
5-7 with a readout time of 0.54 s. No X-ray source is found coincident with
the star’s position.

###  A.42 GL 117

Turnbull (2015) report an age of 1.5 Gyr for GL 117 and Takeda et al. (2007)
derive an upper limit of $<$1.2 Gyr. These ages are in contrast with the very
young age ($\sim$100 Myr) predicted by Mamajek & Hillenbrand (2008). There are
no currently known or candidate exoplanets orbiting GL 117.

GL 117 is robustly detected in XMM-Newton observation 0203060501, and there
are low-level count rate variations observable in the light curve (see Figure
Set 13). The largest of these variations reaches a count rate level that is
$\sim$30% above the median count rate. Given the ample number of X-ray counts
recorded by XMM-Newton, we divide the light curve into sub-exposures but we
find no significant differences in the best fit spectral parameters nor the
predicted luminosities (see Figure Set 14 and Table 6). The bright X-ray
luminosity of the star, particularly when compared to other K dwarfs in our
sample, are consistent with a young age for GL 117.

###  A.43 GL 783

GL 783 is a binary system containing a K3V star (A) and an M4.5V star (B;
Turnbull & Tarter, 2003). The two stars were stated to be separated by
$\sim$8′′ by Allen et al. (2000) but Gaia DR3 positions indicate a separation
of $\sim$2.6′′ (Gaia Collaboration et al., 2023), both of which are
unresolvable by XMM-Newton but potentially spatially resolvable by Chandra
(see Section 3). GL 783B has an apparent visual magnitude of $\sim$12.5 (Allen
et al., 2000), which we use to estimate $L_{\rm bol}$. At the time of this
writing there were no published values of $T_{\rm eff}$ for GL 783B. The
system does not host any currently known or candidate exoplanets.

There is one XMM-Newton observation available for the GL 783 system, and the
individual stars are not resolvable (see Figure 12). The therefore cannot
separate the X-ray emission from the K3V star (GL 783A) from the M4.5V star
(GL 783B) with the currently available observations. The $A^{2}$ statistic
provides evidence that one or both stars may be variable on short timescales
(see Figure Set 13), however there are not a sufficient number of counts to
split the observation into sub-exposures. We attempted to model the spectrum
of the GL 783 system despite the low number of counts; the spectrum required
three thermal plasma components to achieve a statistically acceptable fit (see
Figure Set 14 and Table 6). The coolest thermal plasma component ($\sim$0.1
keV) is similar to the coolest components of the other FGK stars in our
sample, while the temperatures of the hottest ($\sim$1 keV) and intermediate
($\sim$0.3 keV) components more closely resemble components found in the
later-K and M dwarfs in our sample.

Since we cannot separate the X-ray emission of the binary components
individually, we do not use the best-fit spectrum of the GL 783 system as
template for inferring luminosities of other stars in our sample. Follow-up
observations with higher spatial resolution are needed to associate the
observed X-ray emission with one or both binary components.

###  A.44 GL 183

There is some discrepancy in the reported spectral type of GL 183: Gray et al.
(2006) report a spectral type K3V, while K4III is reported by Turnbull (2015).
We adopt a spectral type K3V. There is additional tension in the age
estimates, with a younger age ($\sim$2 Gyr) reported by Mamajek & Hillenbrand
(2008) and older ages reported by Turnbull (2015, 5.3 Gyr) and Takeda et al.
(2007, $<$5.45 Gyr). There are no currently known or candidate exoplanets
orbiting GL 183.

GL 183 was detected in a 18.4 ks XMM-Newton observation (ID number 0780510301)
$\sim$650 net counts. The $A^{2}$ statistic provides evidence that the star
may be variable on short timescales (see Figure Set 13), however there are not
a sufficient number of counts to split the observation into sub-exposures. We
attempted to model the spectrum of GL 183 despite the low number of counts and
found the spectrum required two thermal plasma components to achieve a
statistically acceptable fit (Figure Set 14; see also Table 6), and the fit
parameters were reasonably well-constrained.

###  A.45 GJ 667