the proper motion of the system’s photocenter will deviate from a straight
line unless both components contribute equally to the light output. This
methodology was first used by Bessel (1844) to deduce hidden companions to
Sirius and Procyon. An illustration of the effect, which is known both as
"proper motion anomaly" and as "acceleration", is shown in Figure 3. This
procedure has seen a recent revival now that high quality Hipparcos motions
from the early 1990s and high quality Gaia DR2 motions from the mid-2010s can
be compared.

Figure 3: Schematic diagram demonstrating the concept of proper motion
anomaly. A binary star system, comprised of an A component (solid black orbit
and black points) and a lower-mass B component (dotted pink orbit and pink
points) is shown at four separate times corresponding to approximate start and
end dates of Hipparcos (left pair) and Gaia DR2 (right pair). The center of
mass (grey squares) moves from left to right over time, and the true proper
motion of the system over the Hipparcos and Gaia timeframes is represented by
the two grey arrows. Assuming that the A component dominates the light of the
system, neither Hipparcos nor Gaia will measure this true motion because the
center of light will move with component A as the stars orbit their
barycenter. The black arrow at left thus shows the proper motion that would be
measured by Hipparcos, and the black arrow at right shows the motion measured
by Gaia DR2. The disagreement between these two independent measurements is
termed "proper motion anomaly" and provides evidence that the system has an
unseen component. (For simplicity, we have removed parallactic motion by
showing only those points at the same parallax factor, as depicted by the time
stamps at the bottom of the figure.)

The lack of agreement between the motion measurements is sufficient to
identify a hidden companion, and only a few other measurements are needed to
derive the companion’s mass. This can be computed from the following equation
from Brandt et al. (2019)

$M=\frac{{s^{2}}({a_{PM}}^{2}+{a_{RV}}^{2})^{\frac{3}{2}}}{G({\varomega}a_{PM})^{2}},$
(9)

where $s$ is the projected separation between the companion and host star,
$a_{PM}$ is the host star’s acceleration on the plane of the sky, $a_{RV}$ is
the host star’s acceleration along the line of sight, and $\varomega$ is the
parallax of the system. (See Equation 17 for a different treatment.) This
equation holds only if all measurements can be approximated to refer to the
same orbital epoch. Otherwise, as detailed in Brandt et al. (2019), more
complex orbital fitting is required.

####  4.1.5 Compact objects with gravitational redshifts

Finally, gravitational redshift can be used to measure the mass if the surface
gravity can also be determined. Within the 20-pc sample, this is realistically
measurable only in relatively massive compact objects like white
dwarfs272727Brown dwarfs, like white dwarfs, are electron degenerate but are
less compact and less massive, so their gravitational redshifts are only
$\sim$0.5 km s-1 for the most massive examples. Although this is comparable to
the effect seen for the Sun, $\sim$0.6 km s-1, it is one hundred times smaller
than the effect seen for a 0.8 M⊙ white dwarf ($\sim$50 km s-1).. The observed
velocity shift, $v_{gr}$, due to gravitational redshift is given by

$v_{gr}=\frac{GM}{Rc},$ (10)

where $c$ is the speed of light (e.g., Chandra et al. 2020). Because the
star’s mass is related to its surface gravity, $g$, via the equation

$g=\frac{GM}{R^{2}},$ (11)

the mass can be computed from

$M=\frac{c^{2}{v_{gr}}^{2}}{Gg}.$ (12)

The surface gravity can be measured from the white dwarf’s spectrum by
comparing to model atmospheres. In practice, though, this method cannot be
applied to single white dwarfs because the gravitational redshift is not
separable from the radial velocity. If the star is part of a co-moving
multiple system or is a member of a young cluster or association, however,
then the degeneracy between the radial velocity component and gravitational
redshift component can be broken.

###  4.2 Single objects

Researchers have employed several methods that are capable of measuring the
masses of individual objects. These techniques – gravitational lensing,
asteroseismology, and surface convection monitoring (aka "flickering") – are
described below.

####  4.2.1 Gravitational lensing

Gravitational lensing occurs when a mass moves very close to the line of sight
between an observer and a background object. The mass of the intervening
object acts as a lens that alters the apparent position of the background
source as seen by the observer (Gaudi 2012) and is potentially measurable for
any object. The two temporarily generated images of the background source have
a morphology that is azimuthally asymmetric, and this manifests itself
observationally as a shift in the centroid. The astrometric shift of the
photocenter is given by

${\boldsymbol{\delta}}(t)={\bf u}(t)\frac{\theta_{E}}{u(t)^{2}+2}$ (13)

where $\theta_{E}$ is the angular Einstein radius, which can be expressed as

${\theta_{E}}^{2}=\frac{4GM_{l}}{c^{2}}({D_{l}}^{-1}-{D_{s}}^{-1})$ (14)

(Walker 1995, Lu et al. 2016). Here, $u$ and $\bf{u}$ represent the scalar and
vector time-dependent lens-source separation in the plane of the sky
normalized to $\theta_{E}$, $M_{l}$ is the mass of the lens, and $D_{l}$ and
$D_{s}$ are the distances to the lens and source, respectively. When the
distances to the lens and source are known, the monitoring of the astrometric
shift as a function of time enables a measurement of the mass of the lens.
These equations show that closer lenses produce larger astrometric signals,
which makes this a valuable technique for measuring the masses of nearby
objects, the main limitation being that such encounters of a lens and a
background source happen only rarely and very accurate astrometry is needed to
predict such encounters a priori. This technique has so far been successfully
applied to only two objects in the 20-pc sample (Sahu et al. 2017, Zurlo et
al. 2018) but promises to become more valuable as more accurate Gaia
parallaxes and proper motions become available for stars all across the Milky
Way.

####  4.2.2 Asteroseismology

Asteroseismology is the study of oscillations in stellar atmospheres, and the
characteristics of these oscillations can be used to deduce a star’s physical
parameters. Any star having a mechanism that can drive oscillations – such as
surface convection, pulsations, tidal effects in a close binary, or opacity
effects (the $\kappa$-mechanism) – can potentially have its mass measured.
Equation 52 in Aerts (2021) can be rewritten to show that the stellar mass,
$M$, can be determined from these oscillations using the relation

$M\sim\frac{{\nu_{max}}^{3}{T_{\rm eff}}^{\frac{3}{2}}}{\Delta\nu^{4}}.$ (15)

(In the absence of a definitive theoretical model for convection, the scaling
of this relation is based on observations of the Sun, as described in Kjeldsen
& Bedding 1995.) Here, $\Delta{\nu}$ is the large frequency separation,
$\nu_{max}$ is the frequency of maximum power, and ${T_{\rm eff}}$ is the
effective temperature. The quantity $\nu$ can also be thought of as the
inverse of twice the sound travel time between the stellar center and the
stellar surface (Eq. 39 of Aerts 2021). Figures 4 and 10 of Aerts (2021)
graphically demonstrate how $\nu$ and $\nu_{max}$ are measured in practice.
The effective temperature, ${T_{\rm eff}}$, is obtained by comparing broad-
wavelength spectroscopy of the star to model atmospheres.

####  4.2.3 Surface convection monitoring ("flickering")

The full asteroseismology treatment above requires high-quality data over a
sufficient time baseline with which to resolve the individual oscillation
modes. However, variations in surface convection alone require less exquisite
data and can be used to measure the mass, if certain ancillary quantities have
also been well measured (Stassun et al. 2018). The needed quantity is
$\nu_{max}$ from above, which has been shown to depend on the star’s gravity,
$g$, and effective temperature, $T_{\rm eff}$, through the relation

$g=\frac{\nu_{max}\sqrt{T_{\rm eff}}}{{C}}$ (16)

(Brown et al. 1991), where $C$ is a normalization constant obtained by
calibrating to stars with gravity measurements independently determined via
asteroseismology (Kallinger et al. 2016). The effective temperature is, as
above, obtained by comparing broad-wavelength spectroscopy to model
atmospheres. The mass can then be measured via Equation 11, where the star’s
radius can be measured directly via interferometry or through the Stefan-
Boltzmann Law in Equation 5. The bolometric luminosity can be measured from
the aforementioned broad-wavelength spectrum along with an accurate
trigonometric parallax.

##  5 Multiples Lacking Sufficient Data for Mass Determination

There are some systems for which acceleration has been measured or whose
astrometry indicates the presence of multiple components but for which
insufficient data exist to compute the masses of the individual objects. Such
systems are important to note because their mass accounting is still
incomplete. This serves as an additional source of uncertainty in our mass
function analysis.

###  5.1 Multiples known only through limited acceleration data

Currently, there are many accelerating objects within the 20-pc census that
lack the additional data needed for companion mass computations via Equation
9. We nonetheless still note these as binaries in Table 3.1, and we split out
those cases here for individual discussion.

Khovritchev & Kulikova (2015) have identified likely accelerators by comparing
the proper motion measured between the first and second Palomar Observatory
Sky Surveys (POSS-I and POSS-II; Minkowski & Abell 1963; Reid et al. 1991;
Lasker & STSCI Sky-Survey Team 1998) to a motion derived using first-epoch
data from other sky surveys (2MASS, SDSS, WISE) and their own second-epoch
follow-up astrometry. With these two independent measurements, they can
compare a long-baseline motion over 50 yr to one derived more instantaneously,
over only $\sim$10 yr. Brandt (2021) have similarly intercompared the near-
instantaneous Hipparcos-measured proper motion from the early 1990’s, the
near-instantaneous Gaia-measured motion from the mid- to late-2010’s, and a
long-baseline motion constructed from the Hipparcos-to-Gaia baseline. Both
sets of authors have identified objects with significant motion discrepancies
and labeled these as likely binaries. These objects are noted in Table 3.1
using the column labeled "Accelerator?".

####  5.1.1 Accelerators from POSS vs. recent motion comparison

The Khovritchev & Kulikova (2015) list of $\sim$2400 objects covers only a
portion of the northern sky ($30^{\circ}<{\rm Dec}<70^{\circ}$) for bright
($V<17$ mag), high motion ($\mu>300$ mas yr-1) stars. Within 20 pc of the Sun,
nine such accelerators are identified, only two of which – BD+66 34 and G
96-29 (Capella HL) – were already identified as known multiples in Table 3.1.
The other seven are listed in Table 9.

Table 9: New 20-pc Accelerators from the Khovritchev & Kulikova (2015) Sample J2000 RA & Dec | Name | Sp. Type
---|---|---
(1) | (2) | (3)
00 38 59.04 +30 36 58.4 | Wolf 1056 | M2.5
00 57 02.69 +45 05 09.8 | G 172-30 | M3
01 03 19.84 +62 21 55.8 | Wolf 47 | M5 V
01 38 21.62 +57 13 57.0 | Ross 10 | M2.5
06 01 11.05 +59 35 49.9 | G 192-13 | M3.5 V
19 08 29.93 +32 16 51.6 | G 207-19 | M3.5
23 07 29.92 +68 40 05.2 | G 241-45 | M3

We note that none of these seven objects is identified as a high-significance
accelerator in the Brandt (2021) reference discussed in the following
subsection. This is because the Brandt (2021) Hipparcos-to-Gaia accelerations
could not be computed for these seven stars, as none are in the Hipparcos
Catalog282828Ross 10 has a Hipparcos designation in SIMBAD (HIP 7635) but does
not appear in van Leeuwen (2007).. To further explore the underlying data for
these Khovritchev & Kulikova (2015) accelerators, we have produced finder
charts that show all seven in the POSS-I, POSS-II, 2MASS, and WISE images. A
few of these appear to be blended with a background object at one of the POSS
epochs. The most notably affected are G 172-30, which is blended at POSS-I
with an object fainter by $\Delta$G = 5.8 mag; Wolf 47292929The primary in
this system, BD+61 195, is 295$\arcsec$ away from Wolf 47 itself and would not
be responsible for any acceleration., which is blended at POSS-I with an
object fainter by $\Delta$G = 6.4 mag; and G 192-13, which is blended at POSS-
II with an object fainter by $\Delta$G = 6.2 mag. (Ross 10 moves past a star
of near-equal magnitude in all of the images, the possible blending being
worst at the POSS-II and 2MASS epochs.) This having been noted, whether or not
objects with these magnitude differences could perturb the POSS measurements
enough to affect the 50-yr proper motion measurements is not clear. Future
releases from a longer baseline Gaia data set should determine whether the
accelerations seen for these seven objects are real.

####  5.1.2 Accelerators from Hipparcos vs. Gaia comparisons

The Brandt (2021) list of $\sim$115,000 objects covers the entire sky for
objects in common to Hipparcos and Gaia eDR3 ($G\lesssim 11$ mag). This list
also gives the computed $\chi^{2}$ value between the two proper motions
measured with the best precision, which is usually the Gaia-specific and
Hipparcos-to-Gaia measurements. We conservatively set a false alarm rate of
$Q=e^{-\chi^{2}/2}<0.1\%$, corresponding to $\chi^{2}>13.8$, to select high-
confidence accelerators for analysis here. Using this criterion produces
$\sim$33,750 objects, of which 194 fall within the 20-pc census. These 194 are
denoted in Table 3.1 with a "yes" in the "Accelerator?" column.

Kervella et al. (2022) have also produced a catalog of possible accelerators
based on a comparison of the short-baseline Gaia-specific motions and the
long-baseline Hipparcos-to-Gaia motions. As this list is based on the same
underlying data as the list produced by Brandt (2021), many of the same
accelerators are flagged by both teams. Under the assumption that the
companion mass is much less than that of the primary and that the (circular)
orbit is perpendicular to the line of sight, Kervella et al. (2022) have
further used the proper motion measures to estimate the mass of the hidden
companion using the equation

$m=\Bigl{(}4740.470\frac{\Delta\mu}{\varomega}\Bigr{)}\sqrt{\frac{rM}{G}},$
(17)

where $m$ is the companion mass, $M$ is the primary mass, $G$ is the
gravitational constant, $r$ is the orbital radius, $\Delta\mu$ (the difference
in motion measurements) is in units of mas yr-1, and $\varomega$ is in units
of mas. The constant of 4740.470 is used to convert $\Delta\mu/\varomega$ into
units of m s-1 (Kervella et al. 2019). Companion masses are estimated using
estimated primary masses generally from isochrone fitting for the brightest
stars or from an absolute $K$-band relation for the fainter stars, as
described further in Kervella et al. (2022). Companion masses are dependent
upon the unknown value of the separation between components, so Kervella et
al. (2022) constructed estimates for assumed separations of 3, 5, 10, and 30
AU. In Table 3.1 we include the extrema of these mass estimates in columns
labeled "EstMassAt3AU" and "EstMassAt30AU" for all objects tagged as
accelerators. (In a small number of cases, a Brandt 2021 accelerator was not
deemed to be an accelerator by Kervella et al. 2022, so these estimates are
not given.)

Table 10: 20-pc Accelerators in Known Close Binary/Multiple Systems J2000 RA & Dec | Name | Sp. Type | $\chi^{2}$ |  $M_{est}$ @ 3 AU |  $M_{est}$ @ 30 AU | Note
---|---|---|---|---|---|---
|  |  |  | ($M_{\odot}$) | ($M_{\odot}$) |
(1) | (2) | (3) | (4) | (5) | (6) | (7)
00 49 26.76 $-$23 12 44.9 | HD 4747 AB | G9 V | 12980  | 0.0338 | 0.1532 |  $\cdots$ |
00 50 33.25 +24 49 00.2 | FT Psc AB | M3 V keeaaSome optical spectral types for late-K and M dwarfs include information about chromospheric activity. An "e" generally indicates that H$\alpha$ emission is present. Other values include "(k)" for slight emission reversals or infilling of the Ca II H and K lines, "k" for emission reversals in Ca II H and K that do not rise to the level of the local continuum, "ke" for such emission that rises above the local continuum level, and "kee" for strong emission in Ca II H and K along with H$\beta$ and possibly H$\gamma$ and H$\delta$ (Gray et al. 2003). Because chromospheric activity is variable, these emission-line classification suffixes pertain only to the epoch of spectroscopic observation. | 768.3 | null | null |  $\cdots$ |
00 58 27.94 $-$27 51 25.4 | CD$-$28 302 AB | M3 V | 24.25 | 0.0078 | 0.0772 |  $\cdots$ |
01 41 47.13 +42 36 48.2 | HD 10307 AB | G1 V | 2540  | 0.1971 | 0.6338 |  $\cdots$ |
01 55 57.46 $-$51 36 32.0 |  $\chi$ Eri AB | G9 IV | 569.3 | 0.0970 | 0.1717 |  $\cdots$ |
02 05 04.88 $-$17 36 52.7 | BD$-$18 359 AB | M3 | 6920  | 0.0925 | 0.6897 |  $\cdots$ |
02 10 25.92 $-$50 49 25.5 | HD 13445 AB | K1 V | 78930  | 0.0957 | 0.4335 | 1 known exoplanet |
02 19 10.08 $-$36 46 41.2 | L 440-30 AB | M2.5+ V | 90.10 | 0.0034 | 0.0273 |  $\cdots$ |
02 36 04.90 +06 53 12.4 | HD 16160 AB | K3 V | 6078  | 0.0942 | 0.4769 |  $\cdots$ |
02 45 06.20 $-$18 34 21.4 |  $\tau^{1}$ Eri AB | F6 V | 4366  | 0.1935 | 0.5594 |  $\cdots$ |
02 46 17.28 +11 46 30.9 | HD 17230 AB | K6 V | 60.88 | 0.0022 | 0.0125 |  $\cdots$ |
03 01 51.39 $-$16 35 36.0 | BD$-$17 588 ABC | M2.5 V | 215.2 | 0.0147 | 0.1728 | 2 known exoplanets |
03 48 01.70 +68 40 38.8 | G 221-24 AB | K6 V + M2 V | 1011  | 0.1764 | 1.2275 |  $\cdots$ |
05 08 35.04 $-$18 10 19.4 | L 737-9 AB | M3.5 V | 41.46 | 0.0024 | 0.0169 |  $\cdots$ |
05 19 12.66 $-$03 04 25.7 | HD 34673 AB | K3 V | 2500  | 0.0128 | 0.0739 |  $\cdots$ |
05 22 37.48 +02 36 11.6 | HD 35112 AB | K2.5 V | 192800  | 0.0806 | 1.0007 |  $\cdots$ |
05 28 44.87 $-$65 26 55.2 | AB Dor ACaCb | K2 V k | 1173  | 0.1355 | 0.6137 |  $\cdots$ |
05 32 14.66 +09 49 14.9 | Ross 42 AB | M4 | 17.34 | 0.0010 | 0.0079 |  $\cdots$ |
05 54 22.96 +20 16 34.5 |  $\chi^{1}$ Ori AB | G0 V CH-0.3 | 430.1 | 0.0578 | 0.2097 |  $\cdots$ |
06 10 34.61 $-$21 51 52.7 | Gl 229 AB | M1 V | 13040  | 0.0078 | 0.0544 | 2 known exoplanets |
06 17 16.13 +05 05 59.9 | HD 43587 AaAb | G0 V | 22570  | 0.2482 | 0.7983 |  $\cdots$ |
06 26 10.25 +18 45 24.8 | HD 45088 AaAb | K3 V k | 19680  | 0.0712 | 0.3606 |  $\cdots$ |
06 36 18.29 $-$40 00 23.6 | CD$-$39 2700 AB | K8 V k | 14140  | 0.2114 | 1.4834 |  $\cdots$ |
07 16 19.77 +27 08 33.1 | G 109-55 AB | M2.5 V | 243.4 | 0.0240 | 0.1787 |  $\cdots$ |
07 19 31.27 +32 49 48.3 | BD+33 1505 AB | M0 V | 1301  | 0.0174 | 0.1160 |  $\cdots$ |
07 20 07.37 +21 58 56.3 |  $\delta$ Gem AaAb | F2 V kF0mF0 | 36.94 | 0.0322 | 0.0612 |  $\cdots$ |
07 28 51.36 $-$30 14 49.3 | CD$-$29 4446 AB | M2 | 1377  | 0.1433 | 0.8673 |  $\cdots$ |
07 36 07.07 $-$03 06 38.7 | BD$-$02 2198 AB | M1 V | 18520  | 0.1057 | 0.7205 |  $\cdots$ |
08 31 37.57 +19 23 39.4 | CU Cnc AaAbAc | M4 | 2767  | 0.0466 | 0.3328 |  $\cdots$ |
08 36 25.47 +67 17 41.8 | BD+67 552 AB | M0.5 | 9935  | 0.1336 | 0.8255 |  $\cdots$ |
08 39 07.90 $-$22 39 42.8 | HD 73752 AaAbB | G5 IV | 2804  | 0.2452 | 0.8217 |  $\cdots$ |
08 42 44.53 +09 33 24.1 | BD+10 1857 AaAb | M0 | 186.0 | 0.3145 | 1.9750 |  $\cdots$ |
08 57 04.68 +11 38 48.8 | BD+12 1944 AB | M1.5 | 25170  | 0.0973 | 0.6882 |  $\cdots$ |
09 14 53.65 +04 26 34.2 | HD 79555 AB | K3+ V | 22350  | 0.1772 | 1.0191 | also in Table 13 |
09 29 08.93 $-$02 46 08.2 |  $\tau^{1}$ Hya AaAb | F5.5 IV-V | 2351  | 0.3076 | 0.7993 |  $\cdots$ |
09 32 51.43 +51 40 38.3 |  $\theta$ UMa AB | F5.5 IV-V | 518.5 | 0.0727 | 0.1540 |  $\cdots$ |
09 35 39.50 +35 48 36.5 | 11 LMi AaAb | G8+ V | 496.1 | 0.0091 | 0.0364 |  $\cdots$ |
09 45 40.07 $-$39 02 26.5 | L 462-119 AB | M2.5 V | 1697  | 0.1046 | 0.8348 |  $\cdots$ |
09 53 11.78 $-$03 41 24.4 | BD$-$02 3000 AB | M2 | 80.34 | 0.0052 | 0.0351 |  $\cdots$ |
11 11 33.15 $-$14 59 28.9 | HD 97233 AB | K5 V (k) | 471.7 | 0.1108 | 0.7535 |  $\cdots$ |
11 21 26.67 $-$20 27 14.0 | HD 98712 A | K6 V ke | 6618  | 0.1439 | 0.8376 |  $\cdots$ |
12 00 44.46 $-$10 26 46.1 | HD 104304 AB | G8 IV | 57500  | 0.2325 | 0.8631 |  $\cdots$ |
12 23 33.20 +67 11 18.5 | G 237-64 AB | M2.5 | 1767  | null | null |  $\cdots$ |
12 28 57.59 +08 25 31.1 | Wolf 414 AB | M3.5 V + M5 V | 2059  | 0.1223 | 1.0562 |  $\cdots$ |
12 44 14.55 +51 45 33.4 | HD 110833 AaAbB | K3 | 95.44 | 0.0307 | 0.1438 |  $\cdots$ |
13 00 46.56 +12 22 32.7 | BD+13 2618 AB | M1.5 | 3393  | 0.0530 | 0.3667 | 1 known exoplanet |
13 19 33.59 +35 06 36.6 | BD+35 2436 AaAb | M1 | 79.89 | 0.0151 | 0.0948 |  $\cdots$ |
13 47 15.74 +17 27 24.8 |  $\tau$ Boo AB | F7 IV-V | 7745  | 0.1133 | 0.2944 | 1 known exoplanet |
13 52 35.85 $-$50 55 18.1 | HD 120780 AaAb | K2 V | 479.5 | 0.0332 | 0.1594 |  $\cdots$ |
14 03 32.34 +10 47 12.3 | HD 122742 AB | G6 V | 640.6 | 0.0597 | 0.2446 |  $\cdots$ |
14 54 29.24 +16 06 03.8 | BD+16 2708 ABaBb | M3 V | 2764  | 0.0089 | 0.0654 |  $\cdots$ |
15 41 16.57 +75 59 34.0 | Ross 1057 AB | M3.5 | 1780  | 0.0576 | 0.6208 |  $\cdots$ |
15 44 01.82 +02 30 54.6 |  $\psi$ Ser ABaBb | G5 V | 4965  | null | null |  $\cdots$ |
16 05 40.48 $-$20 27 00.1 | HD 144253 AB | K3 V | 189.3 | 0.0439 | 0.2222 |  $\cdots$ |
16 28 28.14 $-$70 05 03.8 |  $\zeta$ TrA AaAb | F9 V | 47.58 | 0.0110 | 0.0369 |  $\cdots$ |
17 09 31.54 +43 40 52.8 | G 203-47 AB | M3.5 V + wd | 51.01 | 0.0072 | 0.0830 |  $\cdots$ |
17 19 03.84 $-$46 38 10.4 | 41 Ara A | G9 V | 832.7 | 0.0160 | 0.0689 |  $\cdots$ |
17 30 11.20 $-$51 38 13.1 | CD$-$51 10924 AB | M0 V | 270.2 | 0.0044 | 0.0274 | 4 known exoplanets |
17 34 59.62 +61 52 28.2 | 26 Dra AB | G0 IV-V | 11960  | 0.3382 | 1.1333 |  $\cdots$ |
17 46 14.42 $-$32 06 08.4 | CD$-$32 13298 AaAb | M3 V | 27.02 | 0.0080 | 0.0597 |  $\cdots$ |
17 46 27.55 +27 43 14.6 |  $\mu^{1}$ Her AaAb | G5 IV | 12600  | 0.1266 | 0.4130 |  $\cdots$ |
18 07 01.59 +30 33 43.6 | b Her AB | F9 V | 66470  | 0.0973 | 1.0649 |  $\cdots$ |
18 10 26.15 $-$62 00 08.0 |  $\iota$ Pav AB | G0 V | 192.8 | 0.0356 | 0.1192 |  $\cdots$ |
18 57 01.64 +32 54 04.7 | HD 176051 AB | F9 V | 38160  | null | null |  $\cdots$ |
19 23 34.01 +33 13 19.1 | HD 182488 AB | G9+ V | 1700  | 0.0085 | 0.0349 |  $\cdots$ |
19 31 07.97 +58 35 09.6 | HD 184467 AB | K2 V | 27.50 | 0.0314 | 0.1422 |  $\cdots$ |
19 54 17.74 $-$23 56 27.9 | HD 188088 AaAb | K2 IV (k) | 15.46 | 0.0011 | 0.0050 |  $\cdots$ |
20 04 06.22 +17 04 12.7 | 15 Sge AB | G0 V | 31560  | 0.0890 | 0.3162 |  $\cdots$ |
20 05 09.78 +38 28 42.6 | HD 190771 AB | G2 V | 12760  | 0.0639 | 0.2319 |  $\cdots$ |
20 10 19.57 $-$20 29 36.4 | HD 191391 AB | K6 V k | 38130  | 0.0819 | 0.4895 |  $\cdots$ |
20 44 21.95 +19 44 58.7 | HD 352860 AB | M0.5 V | 1583  | 0.1875 | 1.2533 |  $\cdots$ |
20 56 48.54 $-$04 50 49.1 | Ross 193 AaAb | M3 | 72.54 | 0.0050 | 0.0395 |  $\cdots$ |
21 00 05.39 +40 04 12.7 | BD+40 883 AaAbB | M2 | 46070  | 0.1350 | 0.8783 |  $\cdots$ |
21 19 45.63 $-$26 21 10.4 | HD 202940 AaAbB | G7 V | 600.4 | 0.0172 | 0.1989 |  $\cdots$ |
21 49 05.76 $-$72 06 09.1 | CD$-$72 1700 AB | M1 | 19020  | null | null |  $\cdots$ |
22 07 00.67 +25 20 42.4 |  $\iota$ Peg AaAb | F5 V | 66.25 | 0.0205 | 0.0533 |  $\cdots$ |
22 18 15.61 $-$53 37 37.5 | HD 211415 AB | G0 V | 612.3 | 0.0132 | 0.0478 |  $\cdots$ |
22 36 09.69 $-$00 50 29.8 | HD 214100 AB | M1 V | 27200  | 0.0672 | 0.4472 |  $\cdots$ |
22 38 45.57 $-$20 37 16.1 | FK Aqr AaAb | M2 | 47.50 | 0.0014 | 0.0091 |  $\cdots$ |
23 01 51.54 $-$03 50 55.4 | HD 217580 AB | K2.5 V | 232.6 | 0.0591 | 0.2966 |  $\cdots$ |
23 39 20.91 +77 37 56.5 |  $\gamma$ Cep AB | K1 III | 4771  | 0.2376 | 0.4206 | 1 known exoplanet |
23 52 25.41 +75 32 40.4 | HD 223778 AaAbB | K3 V | 124300  | 0.1261 | 0.5909 |  $\cdots$ |
23 55 39.78 $-$06 08 33.4 | BD$-$06 6318 AB | M2.5 V k | 2137  | 0.0497 | 0.3530 |  $\cdots$ |
Table 11: 20-pc Accelerators whose only Close Companions are Known Exoplanets J2000 RA & Dec | Name | Sp. Type | $\chi^{2}$ |  $M_{est}$ @ 3 AU |  $M_{est}$ @ 30 AU | # of | Distance to Stellar
---|---|---|---|---|---|---|---
|  |  |  | ($M_{\odot}$) | ($M_{\odot}$) | Known Exoplanets | Companion (AU)
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8)
00 16 12.68 $-$79 51 04.2 | HD 1237 A | G8.5 V (k) | 26.86 | 0.0019 | 0.0080 | 1 | 70 |
00 18 22.88 +44 01 22.6 | GX And | M1.5 V | 2456  | 0.0013 | 0.0114 | 2 | 122 |
03 32 55.84 $-$09 27 29.7 |  $\epsilon$ Eri | K2 V (k) | 33.89 | 0.0013 | 0.0062 | 1 |  $\cdots$ |
04 52 05.73 +06 28 35.6 | Wolf 1539 | M3.5 | 35.10 | 0.0014 | 0.0137 | 1 |  $\cdots$ |
05 11 40.59 $-$45 01 06.4 | Kapteyn’s Star | sdM1 p | 29.58 | 0.0001 | 0.0007 | 1 |  $\cdots$ |
05 37 09.89 $-$80 28 08.8 |  $\pi$ Men | G0 V | 60.98 | 0.0073 | 0.0244 | 3 |  $\cdots$ |
07 54 10.88 $-$25 18 11.4 | CD$-$24 6144 | M0 | 18.89 | 0.0007 | 0.0043 | 2 | 7110 |
09 14 24.68 +52 41 10.9 | HD 79211 | K7 V | 27.48 | 0.0059 | 0.0401 | 1 | 108 |
10 08 43.14 +34 14 32.1 | HD 87883 | K2.5 V | 713.6 | 0.0087 | 0.0408 | 1 |  $\cdots$ |
16 10 24.32 +43 49 03.5 | 14 Her | K0 IV-V | 1009  | 0.0126 | 0.0502 | 2 |  $\cdots$ |
16 12 41.78 $-$18 52 31.8 | LP 804-27 | M3 V | 331.0 | 0.0076 | 0.0562 | 1 |  $\cdots$ |
20 03 37.41 +29 53 48.5 | HD 190360 | G7 IV-V | 14.83 | 0.0016 | 0.0054 | 2 | 2847 |
21 33 33.98 $-$49 00 32.4 | HD 204961 | M2 | 278.3 | 0.0008 | 0.0062 | 2 |  $\cdots$ |
22 03 21.65 $-$56 47 09.5 |  $\epsilon$ Ind A | K4 V (k) | 287.5 | 0.0030 | 0.0157 | 1 | 1464 |

We divide the resulting list of 20-pc accelerators into three subgroups. The
first, listed in Table 10, comprises eighty-three objects in known close
binary and multiple systems. For all of these, the host star is known to have
a close-in companion that Gaia eDR3 fails to detect or provide a full
astrometric solution for, and these companions range in mass from the
substellar regime into the planetary regime. For a host star at a distance of
10 pc, its Hipparcos-to-Gaia acceleration can be detected if the companion has
a separation below a few$\times$100 AU (Figure 12 from Kervella et al. 2022).
Companions at this separation range can also be detected with high-resolution
imaging techniques or via radial velocity monitoring, and some have
independently measured masses. As one example, the companion in the 19.5-yr
spectroscopic binary HD 10307 AB has a measured dynamical mass from Torres
(2022) of 0.254$\pm$0.019 M⊙, and that system has $a=7.7$ AU, $i=100^{\circ}$
and $e=0.44$. The Kervella et al. (2022) companion mass estimates of 0.20 M⊙
at 3 AU and 0.63 M⊙ at 30 AU bracket the dynamically measured values well, as
the assumptions used were reasonable for this system. As another example, the
companion to the 1.35-yr spectroscopic binary HD 184467 AB has a measured
dynamical mass of 0.868$\pm$0.025 M⊙ (Piccotti et al. 2020), and the system
has $a=0.7$ AU, $i=145^{\circ}$ and $e=0.34$ (Arenou et al. 2000). The
Kervella et al. (2022) companion mass estimate of 0.03 M⊙ at 3 AU compares
unfavorably to the measured value possibly because of the Kervella et al.
(2022) assumption that the secondary mass is much less than that of the
primary. This demonstrates that, although the Kervella et al. (2022) companion
mass estimates listed in Table 10 provide a guide as to whether the companion
causing the acceleration is already known or is a still hidden member,
additional astrometric data is needed before the masses can be reliably
measured. As can be seen from the full entries in Table 3.1, many objects in
Table 10 are triples, so it is also unclear how many objects are contributing
to the measured acceleration.

The second list, shown in Table 11, gives fourteen objects known to host
exoplanets but lacking any "close" stellar or substellar companions. Here we
define "close" to mean within $\sim$50 AU. Six of these objects, as listed in
the final column of the table, have more widely separated stellar companions
at apparent separations of $\gtrsim$70 AU. The Kervella et al. (2022) mass
estimates for all fourteen of these objects are quite low and, for assumed
separations of a few AU, correspond to masses traditionally thought of as
being in the planetary range. Thus, the accelerations for these objects are
likely caused by the known exoplanet(s) in the system. Kervella et al. (2022)
provides additional analysis on the stars $\epsilon$ Eri, Kapteyn’s Star,
$\epsilon$ Ind A, and $\pi$ Men, while noting that Kapteyn’s Star has no
significant proper motion anomaly as measured by them.

The final list, shown in Table 12, has ninety-seven objects whose closest
known companions are resolved by Gaia eDR3 or have no known companions at all.
For many of these, the nearest known companion falls close enough to the
accelerator star ($\lesssim$100 AU; Figure 12 of Kervella et al. 2022) that it
may be the object causing the acceleration. Examples are CD$-$44 3045 A, VV
Lyn Aa, CD$-$36 6589 A, Ross 52 A, BD+45 2247 A, and Wolf 1225 A. Objects for
which the nearest known companion lies beyond this separation or for which no
companions are currently known are the hosts most likely to harbor new
additions to the 20-pc census. Examples of stars with likely hidden companions
are G 32-7, CD$-$22 526, HD 13579, LP 837-53, HD 43162 A, HD 52698, G 250-34,
BD$-$17 3088, $\mu$ Vir, $\beta$ TrA, and $\theta$ Cyg.

Tables 10-12 highlight that the accounting of all components within the 20-pc
census is still incomplete, as there is overwhelming evidence of additional,
tightly separated companions. As only $<$200 of the $\sim$3,000 Gaia-detected
primaries show such evidence, it is tempting to conclude that our tally of
higher mass (non-exoplanet) companions is nearing completion. We caution,
however, that our criteria for selecting accelerators was set very
conservatively and that many real accelerators likely exist with a measured
significance below our cutoff value. As the time baseline of Gaia observations
is extended, accelerations will be increasingly sensitive to longer-period
companions that, for higher (non-exoplanet) masses, are potentially verifiable
with direct imaging techniques. Furthermore, Gaia observations over this same
extended time baseline will remove the need to compare to the shallower
Hipparcos data, enabling acceleration data for lower-mass primaries between
the Hipparcos and Gaia limits ($11\lesssim G\lesssim 21$ mag). Finally, less
than a third of all systems in the 20-pc census of Table 3.1 have both a
Hipparcos entry and a Gaia DR3 astrometric solution, so many objects within
our sample volume are unavailable for similar acceleration analysis.

Table 12: 20-pc Accelerators with More Distant (or No Known) Companions J2000 RA & Dec | Name | Sp. Type | $\chi^{2}$ |  $M_{est}$ @ 3 AU |  $M_{est}$ @ 30 AU | # of | Dist. to Next | Note
---|---|---|---|---|---|---|---|---
|  |  |  | ($M_{\odot}$) | ($M_{\odot}$) | Components | Nearest Known |
|  |  |  |  |  | in System | Member (AU) |
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8) | (9)
00 05 41.02 +45 48 43.6 | HD 38 A | K6 V | 6209  | 0.0219 | 0.1406 | 2 | 70 |  $\cdots$ |
00 16 14.63 +19 51 37.5 | G 32-7 | M4 | 14.42 | 0.0082 | 0.0677 | 3 | 387 |  $\cdots$ |
00 45 48.29 $-$41 54 33.1 | HD 4378 A | K5 | 64.43 | 0.0021 | 0.0120 | 2 | 75 |  $\cdots$ |
00 49 06.29 +57 48 54.6 |  $\eta$ Cas A | F9 V | 659.4 | 0.0220 | 0.0780 | 2 | 79 |  $\cdots$ |
00 49 09.90 +05 23 19.0 | Wolf 28 | DZ7.4 | 1621  | 0.0007 | 0.0044 | 1 |  $\cdots$ |  $\cdots$ |
01 03 14.15 +20 05 52.3 | G 33-35 A | M1.5 | 28.44 | 0.0040 | 0.0300 | 2 | 40 |  $\cdots$ |
01 32 26.20 $-$21 54 18.4 | CD$-$22 526 | M1.5 V (k) | 151.0 | 0.0039 | 0.0284 | 1 |  $\cdots$ |  $\cdots$ |
01 39 47.56 $-$56 11 47.2 | p Eri B | K2 V | 1077  | 0.0184 | 0.0864 | 2 | 94 |  $\cdots$ |
02 15 42.55 +67 40 20.3 | HD 13579 | K2 | 58.54 | 0.0041 | 0.0184 | 2 | 736 |  $\cdots$ |
02 37 52.79 $-$58 45 11.1 | L 174-28 | M3 V | 14.88 | 0.0010 | 0.0111 | 1 |  $\cdots$ |  $\cdots$ |
03 12 04.53 $-$28 59 15.4 |  $\alpha$ For A | F6 V | 2451  | 0.0692 | 0.1861 | 3 | 75 |  $\cdots$ |
03 16 13.83 +58 10 02.5 | Ross 370 A | M2 | 130.2 | 0.0058 | 0.0441 | 2 | 69 |  $\cdots$ |
03 23 35.26 $-$40 04 35.0 | HD 21175 A | K1 V | 69960  | 0.0950 | 0.4497 | 2 | 46 |  $\cdots$ |
03 48 01.03 +68 40 22.4 | HD 23189 | K2 V | 41.23 | 0.0036 | 0.0222 | 3 | 303 |  $\cdots$ |
03 57 28.70 $-$01 09 34.1 | HD 24916 A | K4 V | 62.16 | null | null | 3 | 168 |  $\cdots$ |
04 31 11.51 +58 58 37.5 | G 175-34 | M4.5 V | 208.4 | 0.0018 | 0.0203 | 2 | 57 |  $\cdots$ |
04 53 31.20 $-$55 51 37.1 | CD$-$56 1032 A | M3 V | 117.4 | 0.0032 | 0.0265 | 2 | 83 |  $\cdots$ |
05 03 23.90 +53 07 42.5 | BD+52 911 A | M0.5 | 82.34 | 0.0023 | 0.0167 | 2 | 78 |  $\cdots$ |
05 45 48.28 +62 14 12.4 | BD+62 780 | M0 | 94.82 | 0.0013 | 0.0078 | 1 |  $\cdots$ |  $\cdots$ |
05 55 43.21 $-$26 51 23.4 | LP 837-53 | M2.5 V | 21.48 | 0.0010 | 0.0070 | 1 |  $\cdots$ |  $\cdots$ |
06 13 45.30 $-$23 51 43.0 | HD 43162 A | G6.5 V | 16.04 | 0.0013 | 0.0053 | 4 | 408 |  $\cdots$ |
06 33 43.28 $-$75 37 48.0 | L 32-9 | M3 | 25.61 | 0.0013 | 0.0106 | 2 | 192 |  $\cdots$ |
06 37 11.23 $-$50 02 17.7 | CD$-$49 2340 A | K8 V (k) | 7109  | 0.0262 | 0.1645 | 2 | 35 |  $\cdots$ |
06 57 46.63 $-$44 17 28.2 | CD$-$44 3045 B | M3 | 21590  | 0.0283 | 0.2796 | 2 | 19 |  $\cdots$ |
07 01 13.73 $-$25 56 55.5 | HD 52698 | K1 V (k) | 734.1 | 0.0491 | 0.2224 | 1 |  $\cdots$ | Also in Table 13 |
07 07 50.42 +67 12 04.9 | G 250-34aaA single epoch of Keck/NIRC2 data is available in the Keck Observatory Archive for this object. The $K_{s}$\- and $H$-band observations taken on 22 Jan 2021 (UT; PI: Crepp; Program ID: D313) show an object $0{\farcs}22$ from G 250-34 at a position angle of 328∘ and $\Delta{K_{s}}=2.9$ mag. No other high-resolution observations were found in other archives to help confirm or refute this possible companion. | M1.5 | 16320  | 0.1070 | 0.7770 | 1 |  $\cdots$ |  $\cdots$ |
07 31 57.71 +36 13 10.1 | VV Lyn Aa | M3 | 48470  | 0.0862 | 0.6179 | 3 | 19 |  $\cdots$ |
07 57 57.78 $-$00 48 51.9 | HD 65277 A | K3+ V | 645.3 | 0.0081 | 0.0409 | 2 | 95 |  $\cdots$ |
08 08 13.19 +21 06 18.2 | BD+21 1764 A | K7 V | 140.8 | 0.0058 | 0.0359 | 4 | 190 |  $\cdots$ |
08 50 42.30 +07 51 52.5 | BD+08 2131 A | K5 V | 206.6 | null | null | 2 | 21 |  $\cdots$ |
09 01 17.48 +15 15 56.8 | HD 77175 A | K5 | 210.8 | 0.0062 | 0.0361 | 2 | 93 |  $\cdots$ |
09 06 45.35 $-$08 48 24.6 | BD$-$08 2582 | M0 | 143.1 | 0.0007 | 0.0042 | 2 | 123 |  $\cdots$ |
09 14 22.77 +52 41 11.8 | HD 79210 | M0 V | 178.2 | 0.0053 | 0.0344 | 2 | 108 |  $\cdots$ |
09 43 55.61 +26 58 08.4 | Ross 93 | M3.5 | 23.32 | 0.0023 | 0.0211 | 1 |  $\cdots$ |  $\cdots$ |
10 12 08.15 $-$18 37 04.1 | BD$-$17 3088 | M0 | 267.3 | 0.0561 | 0.3563 | 2 | 7225 |  $\cdots$ |
10 31 24.22 +45 31 33.8 | BD+46 1635 A | K7 V | 111.8 | 0.0026 | 0.0151 | 2 | 74 |  $\cdots$ |
10 41 09.30 $-$36 53 43.7 | CD$-$36 6589 A | M0.5 V | 101600  | 0.1453 | 0.9567 | 2 | 13 |  $\cdots$ |
10 41 51.83 $-$36 38 00.1 | CD$-$35 6662 | M0 V (k) | 31.36 | 0.0045 | 0.0306 | 2 | 305 |  $\cdots$ |
11 05 28.58 +43 31 36.3 | BD+44 2051 A | M1 V | 25.09 | 0.0004 | 0.0040 | 2 | 156 |  $\cdots$ |
11 11 19.48 +43 25 02.4 | G 122-2 A | M2.5 V | 140.0 | 0.0178 | 0.1259 | 2 | 60 |  $\cdots$ |
11 15 11.90 +73 28 30.7 | HD 97584 A | K3 | 47.01 | 0.0014 | 0.0072 | 2 | 94 |  $\cdots$ |
11 34 29.49 $-$32 49 52.8 | 20 Crt A | K0- V | 257.4 | 0.0039 | 0.0175 | 2 | 146 |  $\cdots$ |
11 45 34.44 $-$20 21 12.4 | LP 793-33 | M2.5 V | 618.2 | 0.0245 | 0.1946 | 2 | 295 |  $\cdots$ |
11 45 42.92 $-$64 50 29.5 | LAWD 37 | DQ6.4 | 37.90 | 0.0006 | 0.0040 | 1 |  $\cdots$ |  $\cdots$ |
12 08 24.82 $-$24 43 44.0 |  $\alpha$ Crv A | F1 V | 34.53 | 0.0166 | 0.0431 | 2 | 47 |  $\cdots$ |
12 23 00.16 +64 01 51.0 | Ross 690 | M3 | 125.9 | 0.0039 | 0.0288 | 1 |  $\cdots$ |  $\cdots$ |
12 41 06.49 +15 22 36.0 | HD 110315 A | K4.5 V | 21830  | 0.0394 | 0.2098 | 2 | 32 |  $\cdots$ |
12 41 39.63 $-$01 26 57.9 |  $\gamma$ Vir A | F2 V | 551.7 | 0.1854 | 0.4774 | 2 | 28 |  $\cdots$ |
13 02 20.69 $-$26 47 13.6 | HD 113194 | K5 V (k) | 82.76 | 0.0092 | 0.0528 | 1 |  $\cdots$ | Also in Table 13 |
13 06 15.40 +20 43 45.3 | BD+21 2486 A | K7 | 4329  | 0.0282 | 0.1623 | 3 | 31 |  $\cdots$ |
13 14 15.14 $-$59 06 11.7 | HD 114837 A | F6 V Fe-0.4 | 76.39 | 0.0117 | 0.0327 | 2 | 84 |  $\cdots$ |
13 16 51.05 +17 01 01.8 | HD 115404 A | K2.5 V (k) | 2205  | 0.0089 | 0.0588 | 2 | 85 |  $\cdots$ |
13 20 58.05 +34 16 44.2 | BD+35 2439 | M1.5 | 199.1 | 0.0042 | 0.0301 | 1 |  $\cdots$ |  $\cdots$ |
13 23 32.78 +29 14 15.0 | HD 116495 A | M0 V | 213.5 | 0.0072 | 0.0414 | 2 | 30 |  $\cdots$ |
13 28 21.08 $-$02 21 37.1 | Ross 486 A | M3 V | 20.64 | 0.0015 | 0.0118 | 2 | 112 |  $\cdots$ |
13 32 44.60 +16 48 39.1 | G 63-36 A | M2.5 V | 51.28 | 0.0258 | 0.1893 | 2 | 45 |  $\cdots$ |
13 47 42.16 $-$32 25 48.1 | HD 120036 A | K6.5 V (k) | 173.5 | 0.0071 | 0.0431 | 2 | 135 |  $\cdots$ |
13 49 04.00 +26 58 47.8 | HD 120476 A | K3.5 V | 63800  | 0.0441 | 0.3214 | 2 | 39 |  $\cdots$ |
13 55 02.56 $-$29 05 25.9 | HD 121271 A | M0 V k | 172300  | 0.0536 | 0.3033 | 2 | 38 |  $\cdots$ |
14 19 00.90 $-$25 48 55.5 | HD 125276 A | F9 V Fe-1.5 CH-0.7 | 904.8 | 0.0170 | 0.0616 | 2 | 64 |  $\cdots$ |
14 42 21.58 +66 03 20.8 | G 239-25 A | M2 | 62.92 | 0.0056 | 0.0497 | 2 | 25 |  $\cdots$ |
14 43 03.62 $-$05 39 29.5 |  $\mu$ Vir | F2 V | 23.50 | 0.0384 | 0.0746 | 2 | 793 |  $\cdots$ |
14 51 23.39 +19 06 01.6 |  $\xi$ Boo A | G7 V | 11500  | 0.0689 | 0.2820 | 2 | 38 |  $\cdots$ |
14 53 51.40 +23 33 21.0 | Ross 52 A | M3 V | 6509  | 0.1125 | 0.9852 | 2 | 10 |  $\cdots$ |
14 57 28.00 $-$21 24 55.7 | HD 131977 | K4 V | 354.6 | 0.0046 | 0.0239 | 4 | 145 |  $\cdots$ |
15 00 55.57 +45 25 34.6 | BD+45 2247 A | M0.5 | 80070  | 0.0951 | 0.6269 | 2 | 23 |  $\cdots$ |
15 47 29.10 $-$37 54 58.7 | HD 140901 A | G7 IV-V | 18.43 | 0.0023 | 0.0085 | 2 | 220 |  $\cdots$ |
15 55 08.56 $-$63 25 50.6 |  $\beta$ TrA | F1 V | 31.48 | 0.0163 | 0.0295 | 1 |  $\cdots$ |  $\cdots$ |
16 16 45.31 +67 15 22.5 | EW Dra | M3 | 15.22 | 0.0007 | 0.0052 | 2 | 695 |  $\cdots$ |
16 20 03.51 $-$37 31 44.4 | CD$-$37 10765 A | M2 V | 2809  | 0.0074 | 0.0656 | 2 | 36 |  $\cdots$ |
16 28 52.66 +18 24 50.6 | HD 148653 A | K2 V | 2069  | 0.0163 | 0.0827 | 2 | 46 |  $\cdots$ |
16 55 25.22 $-$08 19 21.3 | Wolf 629 | M3.5 V | 15.84 | 0.0005 | 0.0059 | 5 | 469 |  $\cdots$ |
16 55 38.01 $-$32 04 03.7 | HD 152606 | K8 V k | 1687  | 0.0142 | 0.0874 | 1 |  $\cdots$ |  $\cdots$ |
16 56 48.57 $-$39 05 38.2 | CD$-$38 11343 A | M3 | 550.3 | 0.0184 | 0.1312 | 3 | 49 |  $\cdots$ |
16 57 53.18 +47 22 00.1 | HD 153557 A | K3 V | 875.6 | 0.0102 | 0.0514 | 3 | 91 |  $\cdots$ |
17 15 20.98 $-$26 36 10.2 | 36 Oph B | K0 V | 29.24 | 0.0017 | 0.0078 | 3 | 30 |  $\cdots$ |
17 21 00.37 $-$21 06 46.6 |  $\xi$ Oph A | F2 V | 82.92 | 0.0248 | 0.0646 | 2 | 71 |  $\cdots$ |
17 35 13.62 $-$48 40 51.1 | CD$-$48 11837 A | M1.5 V | 8601  | 0.0174 | 0.1332 | 2 | 45 |  $\cdots$ |
17 46 34.23 $-$57 19 08.6 | L 205-128 | M4 | 37.12 | 0.0001 | 0.0012 | 1 |  $\cdots$ |  $\cdots$ |
17 57 48.50 +04 41 36.1 | Barnard’s Star | M4 V | 74.89 | 0.0000 | 0.0005 | 1 |  $\cdots$ |  $\cdots$ |
18 42 46.70 +59 37 49.4 | HD 173739 | M3 V | 723.0 | 0.0020 | 0.0209 | 2 | 41 |  $\cdots$ |
18 42 46.89 +59 37 36.7 | HD 173740 | M3.5 V | 91.38 | 0.0026 | 0.0309 | 2 | 41 |  $\cdots$ |
18 57 30.59 $-$55 59 30.8 | HD 175224 A | M1 | 412.4 | 0.0066 | 0.0378 | 2 | 33 |  $\cdots$ |
19 36 26.53 +50 13 16.0 |  $\theta$ Cyg | F3+ V | 103.1 | 0.0277 | 0.0648 | 2 | 2150 | Other component in Table 13 |
19 45 49.75 +32 23 13.7 | G 125-30 | M1.5 | 25.44 | 0.0044 | 0.0353 | 1 |  $\cdots$ |  $\cdots$ |
19 46 23.93 +32 01 01.4 | HD 331161 A | M0.5 V | 420.0 | 0.0068 | 0.0472 | 2 | 79 |  $\cdots$ |
20 02 34.16 +15 35 31.5 | HD 190067 A | K0 V Fe-0.9 | 679.7 | 0.0071 | 0.0322 | 2 | 62 |  $\cdots$ |
20 41 51.13 $-$32 26 06.7 | AT Mic A | M4.5 V | 20360  | 0.0891 | 0.5735 | 3 | 21 |  $\cdots$ |
21 02 40.75 +45 53 05.2 | HD 200560 A | K2.5 V | 6003  | 0.0230 | 0.1106 | 2 | 52 |  $\cdots$ |
21 06 53.94 +38 44 57.9 | 61 Cyg A | K5 V | 317.4 | 0.0068 | 0.0386 | 2 | 111 |  $\cdots$ |
21 06 55.26 +38 44 31.4 | 61 Cyg B | K7 V | 16640  | 0.0070 | 0.0431 | 2 | 111 |  $\cdots$ |
21 07 10.38 $-$13 55 22.5 | HD 200968 A | G9.5 V (k) | 4553  | 0.0180 | 0.0816 | 2 | 64 |  $\cdots$ |
21 38 00.39 +27 43 25.4 | BD+27 4120 A | M0.5+ V | 359.2 | 0.0266 | 0.1860 | 2 | 44 |  $\cdots$ |
22 14 31.41 +27 51 18.7 | G 188-49 A | K7.5 | 20.55 | 0.0023 | 0.0157 | 2 | 59 |  $\cdots$ |
22 23 29.13 +32 27 33.9 | Wolf 1225 A | M3.5 | 5174  | 0.1158 | 0.7867 | 2 | 19 |  $\cdots$ |
23 31 52.17 +19 56 14.1 | BD+19 5116 A | M3.5 V | 17570  | 0.0172 | 0.1549 | 2 | 34 |  $\cdots$ |
23 39 37.39 $-$72 43 19.8 | HD 222237 | K3+ V | 700.6 | 0.0046 | 0.0231 | 1 |  $\cdots$ |  $\cdots$ |

###  5.2 Multiples with large RUWE values

The Gaia Renormalized Unit Weight Error (RUWE) is a measure of the goodness of
fit of the single-star astrometric model to the observed astrometry and is
expected to be $\sim$1.0 if the fit is a good representation (Lindegren et al.
2021b). This parameter is pulled from Gaia DR3 and is listed in the "RUWE"
column of Table 3.1. Values significantly higher than unity can indicate
either that the object is an unresolved, physical multiple (Penoyre et al.
2020) or that some other effect is causing the photocenter to deviate from
expectations. Two examples of the latter are a chance alignment with a
marginally resolved background star or single-star variability that confounds
the RUWE renormalization itself (Belokurov et al. 2020). The typically quoted
value for selecting likely binaries using this statistic is RUWE $>$ 1.4
(e.g., Fabricius et al. 2021), although Stassun & Torres (2021) have shown
that values of 1.0 $<$ RUWE $<$ 1.4 are also highly predictive of unresolved
multiplicity. While the RUWE normalization works well across the full
population of Gaia-measured stars, Penoyre et al. (2022b) note that it does a
somewhat less adequate job when a selection of nearby (d $<$ 100 pc; the GCNS
of Gaia Collaboration et al. 2020) stars alone is analyzed. For that reason,
they define a new statistic, which they term the Local Unit Weighted Error
(LUWE), that improves upon RUWE for these nearer objects.

Values of RUWE and LUWE change with each subsequent release of Gaia data, and
there is valuable information contained within the differences. Later Gaia
releases have data (and astrometric solutions) covering a longer timespan, so
for unresolved multiple systems with periods roughly equal to or longer than
the timespans of the data release, the RUWE (or LUWE) values may continue to
run high or even become larger between Gaia DR2 and Gaia eDR3 simply because
the photocentric displacement caused by orbital motion in an unresolved binary
makes the single-object astrometric solution fit less well with an extended
data set. (See Penoyre et al. 2022a for additional discussion.) Conversely,
unresolved binaries with shorter periods should improve and eventually get
full astrometric solutions in the Gaia non-single star lists.

With these observations in mind, Penoyre et al. (2022b) devised a two-part
criterion to select the most likely hidden multiples in the 100-pc sample: (1)
LUWE${}_{eDR3}>2$ and (2) $\Delta{\rm LUWE}\equiv{\rm LUWE}_{eDR3}-{\rm
LUWE}_{DR2}>-{\rm LUWE}_{eDR3}/3$. Within our 20-pc census, 104 objects meet
these criteria, and these are the ones labeled with a "yes" in the Table 3.1
column named "LUWE_binary?". Of these, 73 are known from previous literature
to be binary and were already labeled as such in our census. The other 31,
listed in Table 13, are newly identified multiples. Nine of these are part of
higher-order multiples, as indicated by the notes in the table. For eight of
these systems, Gaia has detections of both the new high-LUWE object and the
other component (sometimes a double itself) with which it has physical
companionship. The ninth system, however, is a new triple system for which
Gaia detects only the new high-LUWE binary G 43-23 but not the common-proper-
motion T dwarf companion, WISEU J100241.49+145914.9, discussed in Section 2.2.
Note that the LUWE criteria from Penoyre et al. (2022b) are meant to be
conservative, so other hidden binaries will exist with LUWE or $\Delta$LUWE
values outside of the bounds noted above.

Table 13: New 20-pc Multiple Systems Identified Through LUWE

J2000 RA | J2000 Dec | Name | Note | Sp. TypeaaSpectral types are taken from Table 3.1. All are optical types except for LP 440-17, which is a near-infrared type.
---|---|---|---|---
(1) | (2) | (3) | (4) | (5)
01 46 29.35 |  $-$53 39 32.6 | 2MASS J01462935$-$5339325 | 1 | M4.5e
04 34 45.33 |  $-$00 26 46.5 | G 82-33 |  | M4 V
07 01 13.73 |  $-$25 56 55.5 | HD 52698bbThis object is also in the accelerator list of Table 12. |  | K1 V (k)
07 08 07.01 |  $-$22 48 47.3 | LP 840-16 |  | M2
07 49 42.14 |  $-$03 20 34.0 | UCAC4 434-042012 | 2 | M3.5 V
08 25 52.82 | +69 02 01.1 | LP 35-347 |  | M5.5 V
10 02 42.45 | +14 59 13.0 | G 43-23 | 3 | M4 V
10 14 53.12 | +21 23 46.0 | G 54-19 |  | M4.5 V
10 39 45.41 |  $-$44 30 37.0 | TYC 7722-1583-1 |  | M3
11 45 34.44 |  $-$20 21 12.4 | LP 793-33 | 4 | M2.5 V
13 02 20.69 |  $-$26 47 13.6 | HD 113194bbThis object is also in the accelerator list of Table 12. |  | K5 V (k)
13 30 40.95 |  $-$20 39 03.7 | UCAC4 347-066233 |  | M4
13 40 08.79 | +43 46 38.0 | Ross 1026 |  | M3.5
13 58 52.21 | +27 52 14.2 | UPM J1358+2752 |  | $\cdots$
14 17 22.10 | +45 25 46.0 | FBS 1415+456 | 5 | M5/6
14 23 43.74 | +14 26 51.4 | LP 440-17 |  | M7 (NIR)
14 24 18.70 |  $-$35 14 32.7 | 2MASSI J1424187$-$351432 |  | M6.5 V
14 28 17.58 | +05 18 45.8 | G 65-53 | 6 | M3.5
15 10 16.82 |  $-$02 41 08.1 | TVLM 868-110639 |  | M9 V
15 18 31.46 | +20 36 28.2 | UCAC4 554-051865 |  | M4.5 V
15 52 06.55 |  $-$33 59 19.0 | UCAC3 113-186615 |  | $\cdots$
17 33 53.18 | +16 55 13.1 | LSPM J1733+1655 |  | M5.5
18 30 39.45 |  $-$03 56 18.9 | UCAC4 431-076686 |  | M4.5 V
18 43 12.51 |  $-$33 22 46.1 | CD$-$33 13497 | 7 | M1
18 53 25.37 | +02 50 48.7 | G 141-46 |  | M2.5
19 36 14.39 | +50 13 10.1 | UCAC3 281-150921 | 8 | M2/3
20 12 59.94 | +01 12 58.3 | 2MASS J20125995+0112584 |  | M6
20 33 36.67 |  $-$21 20 10.1 | 2MASS J20333668$-$2120096 |  | M2
22 10 13.19 |  $-$71 46 06.2 | PM J22102$-$7146 |  | $\cdots$
22 17 18.97 |  $-$08 48 12.3 | Wolf 1561 A | 9 | M4 V
22 24 24.65 |  $-$58 26 13.6 | UCAC3 64-480761 |  | $\cdots$

Note. —

  * •

(1) 0146$-$5339: This is the 35$\farcm$9-distant companion to the F9 dwarf q01
Eri.

  * •

(2) 0749$-$0320: This is the 3$\farcm$9-distant companion to the M3.5 binary
PM J07498$-$0317 AB.

  * •

(3) 1002+1459: This object also has a 15$\farcs$6-distant companion, WISEU
J100241.49+145914.9, announced in this paper (Section 2.2).

  * •

(4) 1145$-$2021: This is a physical system with the M5e star LP 793-34,
15$\farcs$2 distant.

  * •

(5) 1417+4525: This is the 59$\farcs$2-distant companion of the M0 star BD+46
1951.

  * •

(6) 1428+0518: This is a physical system with the M4 star G 65-54, 1$\farcm$0
distant.

  * •

(7) 1843$-$3322: This is a physical system with the M6 star CE 507,
15$\farcs$0 distant.

  * •

(8) 1936+5013: This is the 1$\farcm$9-distant companion to the F3+ dwarf
$\theta$ Cyg, which is listed amongst the accelerators in Table 12.

  * •

(9) 2217$-$0848: This is a physical system with the M5 dwarf binary Wolf 1561
BaBb, 7$\farcs$9 distant.

###  5.3 Multiplicity (and oddities) identified through color-magnitude
diagrams

In Figures 4-9, we show several color-type, color-color, and color-magnitude
diagrams as a final method for identifying unresolved binaries. These diagrams
also illustrate the rich diversity of colors and absolute magnitudes that
objects within the 20-pc census possess.

Figure 4: Various absolute magnitudes plotted against spectral type for the
20-pc census. See text for details.  Figure 5: Absolute $G$-band magnitude
plotted against various colors for the 20-pc census. See text for details. The
spray of mostly black points (i.e., objects with no measured spectral types)
to the left of the main sequence in the $G_{BP}-G$ vs. $M_{G}$ diagram and to
the right of the main sequence in the $G-G_{RP}$ vs. $M_{G}$ diagram
represents components in close binaries near the Gaia resolution limit. The
$G_{BP}$ and $G_{RP}$ magnitudes are calculated from the fluxes in a
3.5$\times$2.1 arcsec2 field, whereas the $G$ magnitudes are calculated from a
profile fit to a much higher-resolution image (section 8 of Evans et al.
2018). For binaries just above the Gaia resolution limit, this means that per-
component BP and RP fluxes will often include light from the other object,
whereas the $G$ flux will not (Halbwachs et al. 2022). This effect pushes such
objects blueward in $G_{BP}-G$ color and redward in $G-G_{RP}$ color, as these
diagrams show.  Figure 6: Various absolute magnitudes plotted against $J-$W2
color for the 20-pc census. See text for details.  Figure 7: Various absolute
magnitudes plotted against $H-$W2 color for the 20-pc census. See text for
details.  Figure 8: Various colors plotted against spectral type for the 20-pc
census. See text for details.  Figure 9: Various color-color diagrams for the
20-pc census. See text for details.

Each plot shows photometry only for those objects believed to be single
components ("# Components this Row" = 1 in Table 3.1) and whose photometry is
uncontaminated ("2MASS_contam?" and/or "WISE_contam?" not equal to "yes" in
Table 3.1). Each object is color coded by its spectral type, as shown by the
color bar in each figure303030This color palette was chosen so that each color
could also be differentiated by readers with deuteranopia, protanopia,
tritanopia, or achromatopsia.. Preference is given to the near-infrared
spectral type if listed; otherwise, the optical spectral type is used. (It
should be noted that for stars of type A through M, near-infrared
classifications are given in Table 3.1 only when no optical type is available,
so this criterion is only relevant for the L, T, and Y dwarfs.) Each object is
plotted as a solid black dot, the center of which is colored if the spectral
type is known; that is, objects lacking a spectral type appear only as black
dots. Furthermore, for plots that involve $J$ or $K$ bands, preference is
given to MKO magnitudes; otherwise 2MASS magnitudes are used.

In Figure 4, only those objects with absolute magnitude uncertainties below
1.0 mag are shown, to keep the plots more legible. In Figure 5, objects are
shown only if their uncertainties in $M_{G}$ are below 1.0 mag and their color
uncertainties are below 0.5 mag. In Figure 6 (or 7), objects are shown only if
their absolute magnitude uncertainties are below 1.0 mag and their $J-$W2 (or
$H-$W2) uncertainties are also below 1.0 mag. In Figure 8, objects are shown
only if the color uncertainty is less than 0.5 mag for Gaia-based color plots
or less than 1.0 mag for all other colors. In Figure 9, points are shown only
if their color uncertainties are generally less than 0.10-0.20 mag.

We have examined each of these diagrams in detail and have identified objects
that fall significantly far from the common loci of main sequence stars or
white dwarfs to warrant special attention. There are several classes of
objects, however, that we do not discuss in this section but address
elsewhere: (1) Stars with bright magnitudes may be problematic and have quoted
uncertainties insufficiently small to capture these problems. Given that these
bright stars are generally well characterized already, we concentrate only on
those not believed to be main sequence stars (category #4 below). (2) L, T,
and Y dwarfs have already been examined in detail via color-type, color-
magnitude, and color-color diagrams in Kirkpatrick et al. (2021a). (3)
Subgiant, giant, and bright giant stars are discussed in Section 6.1.2. (4)
Low-metallicity (subdwarf) stars are discussed in Section 6.2.2. (5) Young
objects are discussed in Section 6.2.1.

We begin with objects whose placement on these diagrams could potentially
highlight a problem with their measured parallaxes. These are all objects that
the Gaia survey is placing within the 20-pc volume for the first time. With
the exception of the object at $11^{h}59^{m}-36^{\circ}34^{\prime}$, all of
these objects have higher than normal Gaia parallax uncertainties as compared
to objects of similar $G$ magnitude. We discuss each of these individually
below:

  * •

Gaia EDR3 4966072879648455296 (0229$-$3606): This object, whose spectral type
has yet to be determined but whose Gaia eDR3/DR3 parallax is $50.66{\pm}0.61$
mas, falls near or just above the main sequence on most color-magnitude
diagrams. Its apparent magnitudes are similar to 20-pc objects of the same
color, so there is no reason to question its inclusion in our census. Its
location on color-magnitude diagrams along with its high DR3 RUWE value of
5.438 indicate possible unresolved binarity.

  * •

Gaia EDR3 3330473222213987072 (0623+1018): This M3 dwarf (see Section A.2) has
a Gaia eDR3/DR3 parallax of $50.80{\pm}1.55$ mas. The derived $M_{G}$ value is
$\sim$8 magnitudes fainter than that expected for an average M3 dwarf, and the
$M_{W2}$ value is $\sim$9 magnitudes fainter. The Gaia parallax value for this
object is clearly in error, so it has been removed from the 20-pc census.

  * •

Gaia EDR3 3460907947316392704 (1159$-$3634): This is an M9.5 dwarf with a Gaia
eDR3/DR3 parallax of $50.10{\pm}0.18$ mas. Its apparent magnitudes fall within
the range of other M9.5 dwarfs within the 20-pc census. In absolute magnitude,
it falls above the main sequence by as much as a magnitude for objects of
similar color, and its Gaia RUWE value is 1.482. This position on color-
magnitude diagrams cannot be explained by binarity alone, but a slightly
larger parallax in tandem would solve the discrepancy. In any event, there is
no reason to exclude this object from Table 3.1.

  * •

Gaia EDR3 6025146733201615616 (1624$-$3212): This object, of unknown type, has
a Gaia eDR3/DR3 parallax of $59.01{\pm}0.12$ mas. Its apparent magnitude falls
in the range expected for objects of similar color within 20 pc. On plots of
absolute magnitude vs. color, however, it appears anomalous. On the $M_{G}$
vs. $G_{BP}-G_{RP}$ plot, it falls 0.4 mag more luminous than objects of
similar color; on the $M_{G}$ vs. $G-J$ plot, it is also more luminous, but by
2.4 mag. Whether these issues indicate a problem with the measured photometry,
the measured astrometry, or both – or whether the object has an unusual
spectrum – is currently unknown. This object is retained in Table 3.1.

  * •

Gaia DR2 4062191480304598656 (AB) (1736$-$2515): This object, also of unknown
type, has a Gaia DR2 parallax of $60.24{\pm}0.83$ mas and is a known double
(Vrijmoet et al. 2022). Gaia DR3 lists two sources near this position, but
neither have a parallax or proper motion measurement. The apparent magnitude
of the DR2 source is at odds with the range expected for objects of similar
color within 20 pc for many combinations of apparent magnitude vs. color, such
as $G$ vs. $G-$W2, $J$ vs. $J-$W2, and $H$ vs. $J-K_{s}$. The object also has
a very small Gaia-measured proper motion of $25.2{\pm}1.1$ mas yr-1 and lies
near the Galactic Center at $l,b=(2{\fdg}1,+3{\fdg}1)$. This object is most
likely a background object with a faulty parallax, so we have removed it from
the 20-pc census.

  * •

Gaia DR2 1795813379365971072 (2151+2328): No spectrum has been acquired of
this object, and it appears to be a very close double in both Gaia DR2 and
eDR3/DR3. However, only one of these components has a parallax measurement in
DR2, and neither one does in DR3. This object lies well below and blueward of
the main sequence on many apparent magnitude vs. color plots such as $J$ vs.
$G-J$, W2 vs. $G_{BP}-G{RP}$, and $H$ vs. $J-$W2. The object also has a small
Gaia DR2 motion of only $45.7{\pm}0.7$ mas yr-1. This is likely a background
source with a bogus parallax, so we have removed it from the 20-pc census.

The rest of our analysis deals with objects that are outliers for various
other reasons. As discussed below, these reasons include possible unresolved
binarity, unusual atmospheric composition, variability corrupting pan-epoch
colors, and suspected typographical errors in published literature values.

  * •

HD 1237 B (0016$-$7951): For its spectral type of M4, this object has
$G_{BP}-G$ and $G_{BP}-G_{RP}$ colors much bluer than expected, while its
$G-G_{RP}$ color is much redder than expected. No other separate photometry of
the B component is given in Table 3.1. Given that the A component is eight
magnitudes brighter in $G$ than the B component and lies only $4{\farcs}0$
away, we suspect a problem with the measured photometry of B that is not
adequately reflected in its quoted uncertainties.

  * •

EGGR 246 (0041$-$2221): This is an oddly blue white dwarf in colors that use
W1 or W2 magnitudes. The object is also blue relative to other white dwarfs in
$J-H$ color, though normal in $G-J$. This carbon-bearing object has a peculiar
spectral type, DQpec9.3, and is believed to have a mixed hydrogen-helium
atmosphere. The known infrared flux deficit is thought to be caused by
absorption by H2 via collisions with neutral He (Giammichele et al. 2012;
Bergeron et al. 1994, 2022).

  * •

LP 941-19 (0213$-$3345): Although this DA4.5 white dwarf has contaminated WISE
photometry, it falls in an odd position on plots based only on Gaia
photometry. Specifically, at its value of $M_{G}$ it falls $\sim$0.5 mag
blueward of the white dwarf locus in $G_{BP}-G$ and $\sim$0.3 mag redward in
$G-G_{RP}$. There is very little literature on this source, and our spectrum
of it (Section A.2) is the first published. It is not yet clear if this
spectrum differs markedly enough from other DA white dwarfs to account for the
color discrepancies or whether the Gaia magnitudes themselves are at fault.

  * •

HD 21209 A (0323$-$4959): The only oddity with this K dwarf is its anomalously
blue W1$-$W2 color. The value in Table 3.1, which is from the WISE All-Sky
Source Catalog, is W1$-$W2 = $-0.17{\pm}0.06$ mag. Although this is the
preferred WISE catalog for sources of this brightness (W1 = 5.56 mag), the
AllWISE Source Catalog gives a very similar color of W1$-$W2 =
$=-0.15{\pm}0.13$ mag. This color may be due to the slightly subsolar
metallicity of the object ([Fe/H]=$-0.44{\pm}0.19$, Soto & Jenkins 2018;
$-0.41{\pm}0.04$, Sousa et al. 2008; $-0.39{\pm}0.02$, Tsantaki et al. 2013).

  * •

HD 23189 (0348+6840): This early-K dwarf is underluminous for its type at
$M_{G}$, $M_{J}$, $M_{H}$, $M_{Ks}$, and $M_{W2}$. When colors formed from
Gaia-based magnitudes are compared to the mean colors of objects of the same
type, it appears normal, whereas the W1$-$W2 color is slightly bluer than
normal. We suspect that the Gray et al. (2003) type of K2 V is a typographical
error, as independent assessments of the type from spectra, colors, and
luminosity considerations (Adams et al. 1935; Bidelman 1985; Mermilliod 1987;
Stassun et al. 2019) suggest a spectral type closer to K7.

  * •

2MASS J05053461+4648017 (0505+4648): This is an M8 dwarf (see Section A.2)
with a Gaia DR3 parallax value of $56.84{\pm}0.60$ mas. The absolute values
calculated with this parallax are similar to those of other known M8 dwarfs in
the census, and a previous parallax of $69.5{\pm}4.7$ mas (Dittmann et al.
2014) also places it within 20 pc. Its location on color-magnitude diagrams
such as $M_{G}$ vs. $G-$W2 along with its high DR3 RUWE value of 4.823
indicate possible unresolved binarity.

  * •

DENIS J071807.3$-$350220 (0718$-$3502): On the $G_{BP}-G_{RP}$ and $G-J$ vs.
various absolute magnitude diagrams, this object appears to be $\sim$0.7 mag
above the locus of other objects of the same color. This is, therefore, likely
a near equal-magnitude double. This object is also flagged as a possible
binary in the Apps Catalog, again based on its position in color-magnitude
diagrams.

  * •

SCR J0818$-$3110 (0818$-$3110): This DZ white dwarf is an outlier on the $G-J$
vs. W1$-$W2 diagram and, in fact, any diagram involving W1$-$W2 color. This
issue has been hinted at previously in Figure 4 of Kawka et al. (2021), which
shows that the best model fit to existing spectra and photometry fails to
match the observed W1$-$W2 color. Although the effect is known, its reason has
apparently not yet been established and may be caused by variability or
missing opacity sources in the models.

  * •

UPM J0901$-$6526 (0901$-$6526): This K5 dwarf is an outlier on all plots
showing spectral type but appears normal on color-color and color-magnitude
plots. We suggest that the published type of K5 (Riaz et al. 2006) results
from a transcription error in the data for this star and that the actual
spectral type is closer to M5.

  * •

APMPM J1251$-$2121 (1250$-$2121): This M6/6.5 dwarf has a Gaia DR3 parallax of
$56.79{\pm}0.19$ mas. Its apparent magnitudes fall within the range expected
for an M6 dwarf within the 20-pc volume, and a previous parallax measurement
of $57.7{\pm}1.7$ mas (Winters et al. 2015) is in agreement with the Gaia one.
On the $M_{G}$ vs. $G_{BP}-G_{RP}$ diagram, it falls $\sim$0.5 mag above the
main sequence, and on the $M_{G}$ vs. $G-$W2 diagram, it falls $\sim$0.7 mag
above. This and the Gaia DR3 RUWE value of 2.888 suggest unresolved binarity.

  * •

HD 113194 (1302$-$2647): Although this K5 dwarf has a Gaia DR3 parallax with a
relatively large uncertainty ($56.94{\pm}0.19$ mas), the earlier Hipparcos
parallax ($56.87{\pm}1.11$ mas) is in agreement. This object also has a high
Gaia DR3 RUWE value, is listed as a high-LUWE object (see Section 5.2), and
shows acceleration (see Section 4.1.4), in addition to falling $\sim$0.5 mag
above the main sequence on the $G$ vs. $G_{BP}-G_{RP}$ diagram. This object is
almost certainly an unresolved binary. It is also considered to be binary in
the Apps Catalog, based on its position on color-magnitude diagrams. The Gaia
DR3 main catalog reports a radial velocity of $-17.56{\pm}7.24$ km s-1 using
seventeen observations over 920 days, along with an amplitude of radial
velocity variations of 60.25 km s-1, further supporting the hypothesis of
binarity. The P-value for radial velocity constancy (rv_chisq_pvalue) is also
0.0.

  * •

2MASSW J1421314+182740 (1421+1827): This M9.5 dwarf has WISE photometry
contaminated by a background source, but it appears unusual in non-WISE colors
as well. Specifically, it has oddly blue $G_{BP}-G$ and $G_{BP}-G_{RP}$ colors
compared to other objects of similar $M_{G}$. However, it appears normal for
its $M_{G}$ in $G-G_{RP}$, $G-J$, $G-H$, and $J-K_{s}$. This may simply
indicate an issue with the $G_{BP}$ magnitude that the formal uncertainty
fails to adequately capture.

  * •

LP 222-65 (1516+3910): This mid-M dwarf lies consistently $\sim$0.6 mag above
the main sequence relative to objects of the same color and spectral type on
color-magnitude diagrams. This is an isolated object with no obvious problems
with its photometry, so we believe this is an unresolved near-equal magnitude
binary.

  * •

UCAC4 554-051865 (1518+2036): This mid-M dwarf is $\sim$0.6 mag more luminous
than objects of similar color on the $M_{G}$ vs. $G_{BP}-G_{RP}$ and $M_{G}$
vs. $G_{BP}-G$ diagrams and has a large RUWE and LUWE value (see Section 5.2).
It is likely an unresolved binary.

  * •

L 339-19 (1640$-$4559): This M3 dwarf shows anomalously red $G-$W3 and $J-$W3
colors for its absolute magnitude, and even more anomalously red $G-$W4 and
$J-$W4 colors. A more careful look at the WISE images shows that the W3
detection is likely real, but the W4 detection likely is not. The W3
photometry from AllWISE (reported in Table 3.1) is 7.10$\pm$0.05 mag and that
from WISE All-Sky is 6.57$\pm$0.04 mag. In $G-$W3 and using the AllWISE value,
the object lies 1.0 mag redward of objects of the same absolute $G$ magnitude;
using the WISE All-Sky value shows the object to lie 0.5 mag redward. Archival
Spitzer/IRAC and Spitzer/MIPS photometry of this object exists in the GLIMPSE
I Spring ’07 Catalog (Benjamin et al. 2003) and MIPSGAL Archive (Carey et al.
2009) at IRSA: ch1 = 7.830$\pm$0.038, ch2 = 7.781$\pm$0.045, ch3 =
7.724$\pm$0.037 (5.8 $\mu$m), ch4 = 7.705$\pm$0.026 (8.0 $\mu$m), and
[24$\mu$m] = 7.14$\pm$0.24 mag. Running these new data points and the
tabulated Table 3.1 photometry through the Virtual Observatory Spectral energy
distribution Analyzer (VOSA313131http://svo2.cab.inta-csic.es/theory/vosa/;
Bayo et al. 2008) suggests not only that the W4 magnitude is in error but that
the W3 magnitude is spuriously bright relative to the bracketing IRAC and MIPS
data points. The spectral energy distribution is otherwise typical of that of
an M3 dwarf. We therefore conclude that there is no infrared excess in this
object.

  * •

UCAC4 317-104829 (1706$-$2643): This DAH white dwarf is normal in Gaia-only
colors, colors formed using Gaia minus near-infrared magnitudes, and colors
formed from $J$, $H$, and $K_{s}$ magnitudes. It is, however, oddly blue in
W1$-$W2. We assume that this anomalous color may be intrinsic to the star and
a result of its strong magnetic field, although it should be cautioned that
this white dwarf is located against a busy region of the Galactic Plane and
may suffer from contamination in its WISE photometry.

  * •

DENIS-P J1733423$-$165449 (1733$-$1654): This L1 dwarf has WISE photometry
that is contaminated by background sources, but it also shows unusual colors
in Gaia-only measurements. Gaia DR3 lists two other point sources within
$2{\farcs}1$ of this object, so its Gaia photometry may be adversely affected
in a way that the formal uncertainties fail to capture.

  * •

LSPM J1733+1655 (1733+1655): The Gaia DR3 parallax of $60.91{\pm}0.48$ mas has
a relatively large uncertainly for its magnitude and is in disagreement with
an earlier published value of $85.40{\pm}3.30$ mas by Dittmann et al. (2014).
This mid-M dwarf is more luminous than objects of similar color by $\sim$1.6
mag on the $G$ vs. $G_{BP}-G$, $G_{BP}-G_{RP}$, $G-G_{RP}$, $G-J$, and $G-$W2
diagrams, if the Gaia DR3 parallax is used. This overluminosity decreases to
$\sim$0.9 mag if the Dittmann et al. (2014) parallax is used instead. This is
a high-RUWE/LUWE object as well (Section 5.2), and so is likely an unresolved
multiple system with problematic Gaia astrometry. Clark et al. (2022) identify
a candidate companion at separation $0{\farcs}14$ and position angle 101∘ at
epoch 2017.3 and again at separation $0{\farcs}36$ and position angle 63∘ at
epoch 2019.7. C. Gelino also finds a single epoch of Keck/NIRC2 data in the
Keck Observatory Archive for LSPM J1733+1655. These are Br$\gamma$ and
$J$-continuum observations taken on 2015 Jul 10 UT (PI: Hansen; Program ID:
U050N2), from which we measure a separation of $0{\farcs}11$ at position angle
248∘. If we assume all three of these measurements refer to the same star and
it is a stationary background object, we derive motions of LSPM J1733+1655 of
$-0.100\arcsec{\rm yr}^{-1}$ in RA and $-0.051\arcsec{\rm yr}^{-1}$ in Dec,
which can be compared to the measured Gaia DR3 values of $-0.135\arcsec{\rm
yr}^{-1}$ in RA and $-0.130\arcsec{\rm yr}^{-1}$ in Dec. The derived magnitude
and direction of motion lead us to conclude that the background star
hypothesis is sound.

  * •

LP 388-55 A (1735+2634): This late-M dwarf is anomalously red, by 0.25 mag, in
$G-G_{RP}$ color but looks normal in the $G_{BP}-G$ color compared to objects
of similar $M_{G}$ magnitude. Curiously, all Gaia-based absolute magnitudes
($M_{G}$, $M_{GBP}$, and $M_{GRP}$) are consistent with the reported spectral
type. The B component is an early-L that is not directly imaged by Gaia but
may nonetheless be subtly affecting the Gaia magnitudes of the A component.

  * •

LP 44-334 A (1840+7240): This primary in a M6.5 dwarf system has a Gaia DR3
parallax ($52.78{\pm}0.09$ mas) with a relatively large uncertainty for its
magnitude, but this value compares favorably to the earlier published value of
$59.3{\pm}2.2$ mas by Lépine et al. (2009). The $G_{BP}-G$ color is too blue
for its $M_{G}$ value, the $G_{BP}-G_{RP}$ color is normal, and the $G-G_{RP}$
color is too red. These issues are likely caused by the nearness of the B
component, only $0{\farcs}8$ away, which is likely corrupting the photometry
of the A component.

  * •

LP 867-15 (1842$-$2328): The colors for this M0 dwarf are more consistent with
an M4 dwarf than with an M0. Pending spectroscopic verification, we assume
that this object has been misclassified.

  * •

SCR J2012$-$5956 (2012$-$5956): This object, a DC9.9 white dwarf, falls below
the white dwarf locus for most colors. It is very blue relative to other white
dwarfs in $J-K_{s}$, $J-H$, and $H-K_{s}$ but looks like other white dwarfs in
colors made with Gaia-only magnitudes. It is somewhat blue in $G-J$, $G-H$,
and $G-K_{s}$ colors. As with EGGR 246 above, the infrared flux deficit is
believed to be caused by H2-He collision-induced absorption (Giammichele et
al. 2012).

  * •

LEHPM 2-783 (2019$-$5816): This M6.5 dwarf is overluminous in all Gaia-based
colors. (Many other colors are nearly degenerate with absolute magnitude or
type in this spectral type range.) On both the $G_{BP}-G_{RP}$ vs. $M_{H}$ and
the $G-$W2 vs. $M_{W2}$ plots, the overluminosity is $\sim$0.7 mag. Ujjwal et
al. (2020) mark this as a possible member of the $\beta$ Pic Moving Group, and
Riaz et al. (2006) note that it is a strong X-ray emitter with strong
H$\alpha$ emission.

  * •

LP 12-90 (2322+7847): This mid-M dwarf lies above the main sequence by
$\sim$0.75 mag on the $M_{G}$ vs. $G_{BP}-G$ plot. On many other plots of
absolute magnitude vs. color, it lies similarly above (and redward of) the
main sequence. This could be another unresolved binary – if confirmed, this
would make its system with HD 220140 AB a quadruple – but the primary in this
system is a young, naked T Tauri star (Makarov et al. 2007), meaning that its
position may be solely due to its youth.

  * •

ZZ Psc (2328+0514): This white dwarf is anomalously red in colors involving
WISE magnitudes – so much so that it falls far from the white dwarf sequence
itself. On a plot of $M_{G}$ vs. W1$-$W2, for example, it lies substantially
redward of both the white dwarf locus and the main sequence. This object, also
known as G 29-38, is known to have a debris disk around it, the first evidence
of which was uncovered by Zuckerman & Becklin (1987). For an update on this
object, see Cunningham et al. (2022).

##  6 Masses from Estimation

Only a small fraction of objects within the 20-pc census has masses measurable
by methods 1 or 2 described in the introduction of Section 4. For the rest, we
must rely on methods 3 and 4 of that section, which depend on comparison to
empirical trends or to theoretical models.

In the first subsection below, we discuss mass measurements for objects not on
the main sequence – namely, white dwarfs, giants/subgiants, and brown dwarfs.
In the second subsection, we summarize mass estimation for main sequence
stars. In the third subsection, we discuss other complications – youth,
subsolar metallicity, and formation scenario – that may need to be considered
when assigning accurate mass estimates for special objects.

###  6.1 Individual objects not on the main sequence

####  6.1.1 White dwarfs

Masses have been measured via one of the methods described in Section 4 for a
handful of white dwarfs in the 20-pc census, but these represent the end-state
masses of the stellar remnants and are not suitable for analysis of the
initial mass function. Rather, what is needed are the initial masses before
evolution off the main sequence. Techniques have been established that use the
final mass of the remnant to estimate the initial mass of the progenitor.

For white dwarfs lacking a direct mass measurement, one can estimate the final
mass of the white dwarf using one of the following two semi-empirical methods.
The first is to use spectroscopic observations of the depth and width of the
hydrogen Balmer, He I, or He II lines to establish, after comparison to
atmospheric models, the $\log{(g)}$ and $T_{\rm eff}$ for each object. Further
comparison of these two parameters to cooling models provides the remnant mass
(e.g., Tremblay et al. 2011; Genest-Beaulieu & Bergeron 2019; Bergeron et al.
1992; Finley et al. 1997). Whereas this first method is applicable only to DA
(hydrogen atmosphere) or DB (helium atmosphere) white dwarfs, an alternate
method can be used both for objects lacking hydrogen lines as well as for
objects lacking spectroscopic observations. In this second method, masses can
still be estimated if an accurate parallax has been measured. Here, absolute
fluxes across as wide a swath of wavelength space as possible are compared to
model atmospheres to provide $\log{(g)}$ and $R$, from which the mass can be
derived from equation 11 (e.g., Bergeron et al. 2019; Tremblay et al. 2019;
Giammichele et al. 2012; Bergeron et al. 2001; Koester et al. 1979).

The next step is to convert this final mass into an initial mass using an
initial-to-final mass relation (e.g., Weidemann 2000). The empirical form of
this relationship has been established using white dwarfs that are members of
open clusters of known age. As described above, spectroscopic observations of
the Balmer lines in these stars can be compared to atmospheric models to
derive $\log{(g)}$ and $T_{\rm eff}$ for each object. A comparison of these
parameters to cooling models provides both the remnant mass as well as the
cooling time since the object left the tip of the asymptotic giant branch. The
known cluster age minus this cooling time gives the main sequence lifetime of
the object, which can then be related back to an initial mass using
theoretical evolutionary isochrones. This same technique can also be applied
to white dwarfs in globular clusters. Due to their much older ages, these
clusters can provide white dwarfs of lower final mass than those available in
young open clusters. Because these globular clusters are much more distant,
their white dwarfs are faint and more difficult to study, so old low-mass
white dwarfs are still not well represented by cluster methods.

This lack of low-mass examples can be partly mitigated by the use of old, wide
binaries for which the second component can be age dated and the separation
between components is large enough that no mass transfer has occurred during
the system’s evolution. Examples are wide subgiant + white dwarf binaries in
which the system can be dated from its more recently evolved member
(Barrientos & Chanamé 2021), wide F/G/K dwarf + white dwarf binaries in which
the age of the main sequence star can be estimated from activity diagnostics
(Catalán et al. 2008, Zhao et al. 2012), and white dwarf pairs in which
comparison of the higher-mass white dwarf to known cluster white dwarfs can
provide an age for the binary, and the difference in the cooling times for the
white dwarf pair gives the main sequence lifetime of the lower-mass white
dwarf (Andrews et al. 2015). Using the results of these methods, the trend of
final mass with initial mass can be fit. As can be seen from figure 9 of
Barrientos & Chanamé (2021), the relation shows considerable scatter at lower
masses, as the age dating methods for individual systems are generally less
robust than those from clusters. The relations we adopt here are the cluster-
based tripartite parameterization found in equations 4-6 of Cummings et al.
(2018) and the quadripartite parameterization found in table 1 of El-Badry et
al. (2018), The former relation is applicable to white dwarfs with
$0.56<M_{\rm final}<1.24M_{\odot}$ ($0.83<M_{\rm initial}<7.20M_{\odot}$), and
the latter relation, which is based on nearby white dwarfs with accurate Gaia
parallaxes, is applicable to white dwarfs with $0.50<M_{\rm
final}<1.37M_{\odot}$ ($0.95<M_{\rm initial}<8.00M_{\odot}$). We further note
that neither the cluster nor field methods have yet extended the initial-to-
final mass relation below final masses of $0.50M_{\odot}$. (As discussed
further below, white dwarfs with final masses below $0.45M_{\odot}$ require
binary interactions, as a single progenitor would imply an age older than the
Universe; Marsh et al. 1995).

Specifically, we apply the following methodology to assign final masses to
white dwarfs in the 20-pc census. First, we use directly measured masses,
whenever such measurements are available. For others, we use final mass
estimates that are based on accurate parallaxes, high S/N spectra, and/or
broad-wavelength data spanning the white dwarf’s spectral energy distribution.
For all other objects, we resort to the Gaia-centric estimates of Gentile
Fusillo et al. (2021) and Gentile Fusillo et al. (2019). These estimates use
only a small fraction of the white dwarf’s spectral energy distribution –
spanning the Gaia optical bandpasses – and thus lead to separate solutions for
hydrogen- vs. helium-atmosphere objects. When our own follow-up has determined
the spectral type of the object, we use this information to break the
degeneracy; otherwise, a hydrogen-atmosphere object is assumed, as noted in
Table 6.1.1.

Table 14: Mass Measurements and Estimates for White Dwarfs in the 20-pc Census

Name | Abbrev. | WD | Final | Final | Ref. | Initial | Initial | Initial | Initial
---|---|---|---|---|---|---|---|---|---
| Coords | NameaaWhite dwarf designations from McCook & Sion (2016). Those in parentheses are "WD" designations found in SIMBAD but not in McCook & Sion (2016). | Mass | Mass |  | Mass | Mass | Mass | Mass
| (J2000) |  | ($M_{\odot}$) | Method |  | (Cummings)bbThe initial mass derived using the initial-to-final mass relation of Cummings et al. (2018). Uncertainties are derived by propagating the listed uncertainty for the final mass and the uncertainties listed for the coefficients of the initial-to-final mass relations in equations 4-6 of Cummings et al. (2018). | (El-Badry)ccThe initial mass derived using the initial-to-final mass relation of El-Badry et al. (2018). One-sigma uncertainties are chosen to match the initial mass 95.4% (2-sigma) envelope shown in figure 3 of El-Badry et al. (2018). | (Adopted)ddThis is our adopted initial mass. When both a Cummings et al. (2018) estimate and an El-Badry et al. (2018) estimate are available, this adopted mass in the unweighted average of those two. | Method
|  |  |  |  |  | ($M_{\odot}$) | ($M_{\odot}$) | ($M_{\odot}$) |
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8) | (9) | (10)
LAWD 96 | 0002$-$4309 | WD 2359$-$434 | 0.78$\pm$0.03 | Spec | 1 | 3.2$\pm$1.5 | 3.4$\pm$0.2 | 3.3$\pm$1.5 | IFMR
LAWD 1 | 0002$-$3413 | WD 0000$-$345 | 0.88$\pm$0.10 | Phot | 1 | 3.8$\pm$1.3 | 4.7$\pm$0.3 | 4.3$\pm$1.3 | IFMR
LP 464-57 | 0007+1230 | WD 0004+122 | 0.57$\pm$0.15 | Phot | 2 | 1.0$\pm$1.9 | 1.7$\pm$0.1 | 1.4$\pm$1.9 | IFMR
EGGR 381 | 0012+5025 | WD 0009+501 | 0.73$\pm$0.04 | Phot | 1 | 2.9$\pm$1.4 | 3.1$\pm$0.2 | 3.0$\pm$1.5 | IFMR
L 50-73 | 0013$-$7149 | WD 0011$-$721 | 0.59$\pm$0.00 | Phot | 1 | 1.3$\pm$0.5 | 1.9$\pm$0.1 | 1.6$\pm$0.5 | IFMR
G 158-45 | 0014$-$1311 | WD 0011$-$134 | 0.72$\pm$0.07 | Phot | 1 | 2.9$\pm$1.5 | 3.0$\pm$0.2 | 3.0$\pm$1.5 | IFMR
EGGR 246 | 0041$-$2221 | WD 0038$-$226 | 0.53$\pm$0.01 | Phot | 1 |  $\cdots$ | 1.3$\pm$0.1 | 1.3$\pm$0.1 | IFMR
Wolf 28 | 0049+0523 | (WD 0046+051) | 0.68$\pm$0.02 | Phot | 1 | 2.4$\pm$0.7 | 2.8$\pm$0.2 | 2.6$\pm$0.7 | IFMR
Wolf 1516 | 0118+1610 | WD 0115+159 | 0.69$\pm$0.04 | Phot | 1 | 2.5$\pm$0.8 | 2.9$\pm$0.2 | 2.7$\pm$0.8 | IFMR
LP 991-16 A | 0124$-$4240 | WD 0121$-$429 | 0.41$\pm$0.01eeQuoted mass estimate assumes the system is a single white dwarf. | Phot | 1 |  $\cdots$ |  $\cdots$ | 1.9$\pm$0.9 | Ultra-low
LP 991-16 B | 0124$-$4240 |  $\cdots$ |  $\cdots$ |  $\cdots$ | 1 |  $\cdots$ |  $\cdots$ | 1.9$\pm$0.9 | Ultra-low
EGGR 307 | 0125$-$2600 | WD 0123$-$262 | 0.58$\pm$0.00 | Phot | 1 | 1.1$\pm$0.4 | 1.8$\pm$0.1 | 1.5$\pm$0.4 | IFMR
LAWD 10 A | 0137$-$0459 | WD 0135$-$052.1 | 0.47 | Orbit | 4 |  $\cdots$ |  $\cdots$ | 1.0$\pm$0.1 | Low
LAWD 10 B | 0137$-$0459 | WD 0135$-$052.2 | 0.52 | Orbit | 4 |  $\cdots$ | 1.2$\pm$0.1 | 1.2$\pm$0.1 | IFMR
L 88-59 | 0143$-$6718 | WD 0141$-$675 | 0.48$\pm$0.06 | Spec | 1 |  $\cdots$ |  $\cdots$ | 1.0$\pm$0.1 | Low
EGGR 268 | 0151+6425 | WD 0148+641 | 0.81$\pm$0.01 | GravRed | 5 | 3.3$\pm$1.5 | 3.5$\pm$0.2 | 3.4$\pm$1.5 | IFMR
GD 279 | 0152+4700 | WD 0148+467 | 0.63$\pm$0.03 | Spec | 1 | 1.8$\pm$0.6 | 2.3$\pm$0.2 | 2.1$\pm$0.7 | IFMR
$\chi$ Eri B | 0155$-$5136 |  $\cdots$ |  $\cdots$ |  $\cdots$ | - |  $\cdots$ |  $\cdots$ | 1.9$\pm$0.9 | Conjecture
HD 13445 B | 0210$-$5049 | WD 0208$-$510 | 0.597$\pm$0.010 | Accel | 6 | 1.4$\pm$0.5 | 2.0$\pm$0.1 | 1.7$\pm$0.5 | IFMR
EGGR 168 | 0211+3955 | WD 0208+396 | 0.48$\pm$0.10 | Spec | 7 |  $\cdots$ |  $\cdots$ | 1.0$\pm$0.1 | Low
LP 649-66 | 0212$-$0804 |  $\cdots$ | 0.53$\pm$0.02 | Gaia H-atm* | 8 |  $\cdots$ | 1.3$\pm$0.1 | 1.3$\pm$0.1 | IFMR
LP 941-19 | 0213$-$3345 |  $\cdots$ | 0.37$\pm$0.05 | Gaia H-atm* | 8 |  $\cdots$ |  $\cdots$ | 1.9$\pm$0.9 | Ultra-low
EGGR 471 | 0232$-$1411 | WD 0230$-$144 | 0.66$\pm$0.06 | Phot | 1 | 2.1$\pm$0.9 | 2.7$\pm$0.2 | 2.4$\pm$1.0 | IFMR
LP 830-14 | 0235$-$2400 | WD 0233$-$242 | 0.58$\pm$0.00 | Phot | 1 | 1.1$\pm$0.4 | 1.8$\pm$0.1 | 1.5$\pm$0.4 | IFMR
EGGR 473 | 0248+5423 | WD 0245+541 | 0.73$\pm$0.03 | Phot | 1 | 2.9$\pm$1.4 | 3.1$\pm$0.2 | 3.0$\pm$1.5 | IFMR
CPD$-$69 177 | 0310$-$6836 | WD 0310$-$688 | 0.67$\pm$0.03 | Spec | 1 | 2.3$\pm$0.7 | 2.8$\pm$0.2 | 2.5$\pm$0.7 | IFMR
$\alpha$ For Bb | 0312$-$2859 |  $\cdots$ |  $\cdots$ |  $\cdots$ | - |  $\cdots$ |  $\cdots$ | 1.9$\pm$0.9 | Conjecture
EGGR 566 | 0325$-$0149 | WD 0322$-$019 | 0.63$\pm$0.05 | Phot | 1 | 1.8$\pm$0.8 | 2.3$\pm$0.2 | 2.1$\pm$0.8 | IFMR
Wolf 219 | 0344+1826 | WD 0341+182 | 0.57$\pm$0.06 | Phot | 1 | 1.0$\pm$0.9 | 1.7$\pm$0.1 | 1.4$\pm$0.9 | IFMR
G 7-16 | 0400+0814 | WD 0357+081 | 0.61$\pm$0.06 | Phot | 1 | 1.5$\pm$0.9 | 2.1$\pm$0.2 | 1.8$\pm$0.9 | IFMR
${\rm o^{2}}$ Eri B | 0415$-$0739 | (WD 0413$-$077) | 0.573$\pm$0.018 | Orbit | 9 | 1.1$\pm$0.5 | 1.7$\pm$0.1 | 1.4$\pm$0.5 | IFMR
$\epsilon$ Ret B | 0416$-$5917 | WD 0415$-$594 | 0.60$\pm$0.12 | Spec | 7 | 1.4$\pm$1.6 | 2.0$\pm$0.2 | 1.7$\pm$1.6 | IFMR
EGGR 169 | 0425+1211 | WD 0423+120 | 0.65$\pm$0.04 | Phot | 1 | 2.0$\pm$0.7 | 2.6$\pm$0.2 | 2.3$\pm$0.8 | IFMR
EGGR 180 | 0431+5858 | WD 0426+588 | 0.675$\pm$0.051 | Lensing | 10 | 2.3$\pm$0.9 | 2.8$\pm$0.2 | 2.6$\pm$0.9 | IFMR
HD 283750 B | 0436+2709 | WD 0433+270 | 1.12$\pm$0.01 | GravRed | 5 | 6.1$\pm$1.2 | 6.5$\pm$0.4 | 6.3$\pm$1.2 | IFMR
EGGR 41 | 0437$-$0849 | WD 0435$-$088 | 0.53$\pm$0.02 | Phot | 1 |  $\cdots$ | 1.3$\pm$0.1 | 1.3$\pm$0.1 | IFMR
LP 777-1 | 0505$-$1722 | WD 0503$-$174 | 0.53$\pm$0.01 | Spec | 23 |  $\cdots$ | 1.3$\pm$0.1 | 1.3$\pm$0.1 | IFMR
V371 Ori B | 0533+0156 |  $\cdots$ | 0.63$\pm$0.17 | Dyn+mod | 18 | 1.8$\pm$2.2 | 2.3$\pm$0.2 | 2.1$\pm$2.2 | IFMR
EGGR 248 | 0551$-$0010 | WD 0548$-$001 | 0.69$\pm$0.03 | Phot | 1 | 2.5$\pm$0.7 | 2.9$\pm$0.2 | 2.7$\pm$0.8 | IFMR
UCAC4 398-010797 | 0554$-$1035 |  $\cdots$ | 0.68$\pm$0.01 | Gaia He-atm | 8 | 2.4$\pm$0.6 | 2.8$\pm$0.2 | 2.6$\pm$0.7 | IFMR
EGGR 45 | 0555$-$0410 | WD 0552$-$041 | 0.82$\pm$0.01 | Phot | 1 | 3.4$\pm$1.5 | 3.7$\pm$0.2 | 3.6$\pm$1.6 | IFMR
EGGR 290 | 0556+0521 | WD 0553+053 | 0.72$\pm$0.03 | Phot | 1 | 2.9$\pm$1.4 | 3.0$\pm$0.2 | 3.0$\pm$1.4 | IFMR
G 249-36 B | 0605+6049 |  $\cdots$ | 1.03$\pm$0.08 | Dyn+mod | 21 | 5.2$\pm$1.3 | 5.9$\pm$0.4 | 5.6$\pm$1.4 | IFMR
$\alpha$ CMa B | 0645$-$1643 | WD 0642$-$166 | 1.017$\pm$0.025 | GravRed | 11 | 5.1$\pm$1.1 | 5.9$\pm$0.4 | 5.5$\pm$1.1 | IFMR
EGGR 484 | 0647+0231 | WD 0644+025 | 0.85$\pm$0.15 | Spec | 7 | 3.6$\pm$1.8 | 4.2$\pm$0.3 | 3.9$\pm$1.8 | IFMR
LAWD 23 | 0647+3730 | WD 0644+375 | 0.69$\pm$0.03 | Spec | 1 | 2.5$\pm$0.7 | 2.9$\pm$0.2 | 2.7$\pm$0.8 | IFMR
L 454-9 | 0657$-$3909 | WD 0655$-$390 | 0.59$\pm$0.00 | Phot | 1 | 1.3$\pm$0.5 | 1.9$\pm$0.1 | 1.6$\pm$0.5 | IFMR
EGGR 485 | 0700+3157 | WD 0657+320 | 0.60$\pm$0.02 | Phot | 1 | 1.4$\pm$0.5 | 2.0$\pm$0.2 | 1.7$\pm$0.6 | IFMR
SCR J0708$-$6706 | 0708$-$6706 | WD 0708$-$670 | 0.57$\pm$0.00 | Phot | 1 | 1.0$\pm$0.4 | 1.7$\pm$0.1 | 1.4$\pm$0.4 | IFMR
EGGR 52 A | 0730+4810 | WD 0727+482A | 0.51$\pm$0.01 | Phot | 1 |  $\cdots$ | 1.1$\pm$0.1 | 1.1$\pm$0.1 | IFMR
EGGR 52 B | 0730+4810 | WD 0727+482B | 0.65$\pm$0.01 | Phot | 1 | 2.0$\pm$0.6 | 2.6$\pm$0.2 | 2.3$\pm$0.6 | IFMR
EGGR 321 | 0733+6409 | WD 0728+642 | 0.58$\pm$0.00 | Phot | 1 | 1.1$\pm$0.4 | 1.8$\pm$0.1 | 1.5$\pm$0.4 | IFMR
$\alpha$ CMi B | 0739+0513 | WD 0736+053 | 0.553$\pm$0.022 | Orbit | 13 |  $\cdots$ | 1.5$\pm$0.1 | 1.5$\pm$0.1 | IFMR
LAWD 25 | 0740$-$1724 | WD 0738$-$172 | 1.11$\pm$0.05 | GravRed | 5 | 6.0$\pm$1.2 | 6.4$\pm$0.4 | 6.2$\pm$1.3 | IFMR
VB 3 | 0745$-$3355 | (WD 0743$-$336) | 0.55$\pm$0.01 | Phot | 1 |  $\cdots$ | 1.5$\pm$0.1 | 1.5$\pm$0.1 | IFMR
EGGR 426 | 0750+0711 | WD 0747+073.1 | 0.48$\pm$0.01 | Phot | 1 |  $\cdots$ |  $\cdots$ | 1.0$\pm$0.1 | Low
EGGR 427 | 0750+0711 | WD 0747+073.2 | 0.56$\pm$0.01 | Phot | 1 | 0.9$\pm$0.4 | 1.6$\pm$0.1 | 1.2$\pm$0.4 | IFMR
LAWD 26 | 0753$-$6747 | WD 0752$-$676 | 0.73$\pm$0.06 | Phot | 1 | 2.9$\pm$1.5 | 3.1$\pm$0.2 | 3.0$\pm$1.5 | IFMR
SCR J075$3$-2524 | 0753$-$2524 | WD 0751$-$252 | 0.52$\pm$0.01 | Gaia-H | 8 |  $\cdots$ | 1.2$\pm$0.1 | 1.2$\pm$0.1 | IFMR
L 97-3 A | 0806$-$6618 | WD 0806$-$661 | 0.58$\pm$0.03 | Phot | 1 | 1.1$\pm$0.6 | 1.8$\pm$0.1 | 1.5$\pm$0.6 | IFMR
UPM J0812$-$3529 | 0812$-$3529 | (WD 0810$-$353) | 0.70$\pm$0.01 | Gaia H-atm* | 8 | 2.6$\pm$0.7 | 2.9$\pm$0.2 | 2.8$\pm$0.7 | IFMR
G 111-64 | 0814+4845 | WD 0810+489 | 0.57$\pm$0.00 | Phot | 1 | 1.0$\pm$0.4 | 1.7$\pm$0.1 | 1.4$\pm$0.4 | IFMR
SCR J0818$-$3110 | 0818$-$3110 | WD 0816$-$310 | 0.57$\pm$0.00 | Phot | 1 | 1.0$\pm$0.4 | 1.7$\pm$0.1 | 1.4$\pm$0.4 | IFMR
SCR J0821$-$6703 | 0821$-$6703 | WD 0821$-$669 | 0.66$\pm$0.01 | Phot | 1 | 2.1$\pm$0.6 | 2.7$\pm$0.2 | 2.4$\pm$0.6 | IFMR
CD$-$32 5613 | 0841$-$3256 | WD 0839$-$327 | 0.45$\pm$0.05 | Spec | 7 |  $\cdots$ |  $\cdots$ | 1.0$\pm$0.1 | Low
LP 726-1 | 0842$-$1347 | WD 0840$-$136 | 0.57$\pm$0.00 | Phot | 1 | 1.0$\pm$0.4 | 1.7$\pm$0.1 | 1.4$\pm$0.4 | IFMR
$\iota$ UMa Ab | 0859+4802 |  $\cdots$ | 1.00$\pm$0.30 | Dyn+mag | 19 | 4.9$\pm$3.0 | 5.8$\pm$0.4 | 5.3$\pm$3.0 | IFMR
LP 606-32 | 0859$-$0058 | WD 0856$-$007 | 0.52$\pm$0.01 | Gaia He-atm | 8 |  $\cdots$ | 1.2$\pm$0.1 | 1.2$\pm$0.1 | IFMR
EGGR 250 | 0915+5325 | WD 0912+536 | 0.75$\pm$0.02 | Phot | 1 | 3.0$\pm$1.5 | 3.2$\pm$0.2 | 3.1$\pm$1.5 | IFMR
EGGR 252 | 1001+1441 | WD 0959+149 | 0.70$\pm$0.01 | Gaia He-atm | 8 | 2.6$\pm$0.7 | 2.9$\pm$0.2 | 2.8$\pm$0.7 | IFMR
LP 315-42 | 1011+2845 | WD 1008+290 | 0.68$\pm$0.01 | Phot | 1 | 2.4$\pm$0.6 | 2.8$\pm$0.2 | 2.6$\pm$0.7 | IFMR
WT 1759 | 1012$-$1843 | WD 1009$-$184 | 0.59$\pm$0.02 | Phot | 1 | 1.3$\pm$0.5 | 1.9$\pm$0.1 | 1.6$\pm$0.5 | IFMR
EGGR 350 | 1023+6327 | WD 1019+637 | 0.57$\pm$0.05 | Phot | 1 | 1.0$\pm$0.8 | 1.7$\pm$0.1 | 1.4$\pm$0.8 | IFMR
LP 37-186 | 1037+7110 | WD 1033+714 | 0.58$\pm$0.00 | Phot | 1 | 1.1$\pm$0.4 | 1.8$\pm$0.1 | 1.5$\pm$0.4 | IFMR
EGGR 535 | 1038$-$2040 | WD 1036$-$204 | 0.60$\pm$0.01 | Phot | 1 | 1.4$\pm$0.5 | 2.0$\pm$0.2 | 1.7$\pm$0.5 | IFMR
BD$-$18 3019 B | 1045$-$1906 | WD 1043$-$188 | 0.53$\pm$0.11 | Phot | 1 |  $\cdots$ | 1.3$\pm$0.1 | 1.3$\pm$0.1 | IFMR
LAWD 34 | 1057$-$0731 | WD 1055$-$072 | 0.85$\pm$0.04 | Phot | 1 | 3.6$\pm$1.6 | 4.2$\pm$0.3 | 3.9$\pm$1.6 | IFMR
$\xi$ UMa Bb | 1118+3131 |  $\cdots$ |  $\cdots$ |  $\cdots$ | - |  $\cdots$ |  $\cdots$ | 1.9$\pm$0.9 | Conjecture
SCR J1118$-$4721 | 1118$-$4721 | WD 1116$-$470 | 0.57$\pm$0.00 | Phot | 1 | 1.0$\pm$0.4 | 1.7$\pm$0.1 | 1.4$\pm$0.4 | IFMR
Ross 627 | 1124+2121 | WD 1121+216 | 0.61$\pm$0.11 | Spec | 7 | 1.5$\pm$1.5 | 2.1$\pm$0.2 | 1.8$\pm$1.5 | IFMR
20 Crt B | 1134$-$3250 | WD 1132$-$325 | 0.60$\pm$0.01 | Gaia He-atm | 8 | 1.4$\pm$0.5 | 2.0$\pm$0.2 | 1.7$\pm$0.5 | IFMR
GD 140 | 1137+2947 | WD 1134+300 | 0.97$\pm$0.03 | Spec | 1 | 4.7$\pm$1.0 | 5.6$\pm$0.4 | 5.1$\pm$1.1 | IFMR
LAWD 37 | 1145$-$6450 | (WD 1142$-$645) | 0.61$\pm$0.01 | Phot | 1 | 1.5$\pm$0.5 | 2.1$\pm$0.2 | 1.8$\pm$0.5 | IFMR
SSSPM J1148$-$7458 | 1147$-$7457 |  $\cdots$ | 0.488$\pm$0.003 | Gaia He-atm | 8 |  $\cdots$ |  $\cdots$ | 1.0$\pm$0.1 | Low
SDSS J115052.32+683116.1 | 1150+6831 | WD 1148+687 | 0.69$\pm$0.04 | Spec | 2 | 2.5$\pm$0.8 | 2.9$\pm$0.2 | 2.7$\pm$0.8 | IFMR
LP 852-7 | 1205$-$2333 | WD 1202$-$232 | 0.59$\pm$0.03 | Spec | 1 | 1.3$\pm$0.6 | 1.9$\pm$0.1 | 1.6$\pm$0.6 | IFMR
G 197-47 | 1211+5724 | WD 1208+576 | 0.56$\pm$0.09 | Phot | 1 | 0.9$\pm$1.2 | 1.6$\pm$0.1 | 1.2$\pm$1.2 | IFMR
WG 21 | 1226$-$6612 | WD 1223$-$659 | 0.45$\pm$0.02 | Spec | 1 |  $\cdots$ |  $\cdots$ | 1.0$\pm$0.1 | Low
WG 22 | 1238$-$4948 | WD 1236$-$495 | 1.13$\pm$0.14 | Spec | 7 | 6.2$\pm$1.8 | 6.5$\pm$0.4 | 6.3$\pm$1.8 | IFMR
Wolf 457 | 1300+0328 | WD 1257+037 | 0.70$\pm$0.06 | Phot | 1 | 2.6$\pm$1.0 | 2.9$\pm$0.2 | 2.8$\pm$1.0 | IFMR
EGGR 436 | 1308+8502 | WD 1309+853 | 0.71$\pm$0.02 | Phot | 1 | 2.8$\pm$0.7 | 3.0$\pm$0.2 | 2.9$\pm$0.7 | IFMR
ER 8 | 1312$-$4728 | WD 1310$-$472 | 0.63$\pm$0.04 | Phot | 1 | 1.8$\pm$0.7 | 2.3$\pm$0.2 | 2.1$\pm$0.7 | IFMR
LP 854-50 | 1319$-$2147 | WD 1316$-$215 | 0.99$\pm$0.01 | Gaia H-atm | 8 | 4.9$\pm$1.0 | 5.7$\pm$0.4 | 5.3$\pm$1.1 | IFMR
LAWD 45 | 1319$-$7823 | WD 1315$-$781 | 0.69$\pm$0.02 | Phot | 1 | 2.5$\pm$0.7 | 2.9$\pm$0.2 | 2.7$\pm$0.7 | IFMR
BD-07 3632 | 1330$-$0834 | WD 1327$-$083 | 0.50$\pm$0.06 | GravRed | 5 |  $\cdots$ | 1.0$\pm$0.1 | 1.0$\pm$0.1 | IFMR
Wolf 489 | 1336+0340 | (WD 1334+039) | 0.54$\pm$0.03 | Phot | 1 |  $\cdots$ | 1.4$\pm$0.1 | 1.4$\pm$0.1 | IFMR
LSPM J1341+0500 | 1341+0500 | (WD 1338+052) | 0.58$\pm$0.15 | Phot | 2 | 1.1$\pm$1.9 | 1.8$\pm$0.1 | 1.5$\pm$1.9 | IFMR
EGGR 438 | 1348+2334 | WD 1345+238 | 0.45$\pm$0.02 | Phot | 1 |  $\cdots$ |  $\cdots$ | 1.0$\pm$0.1 | Low
PG 1350$-$090 | 1353$-$0916 | WD 1350$-$090 | 0.68$\pm$0.03 | Spec | 1 | 2.4$\pm$0.7 | 2.8$\pm$0.2 | 2.6$\pm$0.7 | IFMR
VVV J141159.32$-$592045.7 | 1411$-$5920 |  $\cdots$ | 0.66$\pm$0.01 | Gaia H-atm* | 8 | 2.1$\pm$0.6 | 2.7$\pm$0.2 | 2.4$\pm$0.6 | IFMR
LP 801-9 | 1447$-$1742 | WD 1444$-$174 | 0.82$\pm$0.05 | Phot | 1 | 3.4$\pm$1.6 | 3.7$\pm$0.2 | 3.6$\pm$1.6 | IFMR
G 137-24 | 1535+1247 | WD 1532+129 | 0.57$\pm$0.15 | Phot | 2 | 1.0$\pm$1.9 | 1.7$\pm$0.1 | 1.4$\pm$1.9 | IFMR
HD 140901 B | 1547$-$3755 | WD 1544$-$377 | 0.58$\pm$0.01 | GravRed | 5 | 1.1$\pm$0.5 | 1.8$\pm$0.1 | 1.5$\pm$0.5 | IFMR
CD$-$38 10980 | 1623$-$3913 | WD 1620$-$391 | 0.65$\pm$0.01 | GravRed | 5 | 2.0$\pm$0.6 | 2.6$\pm$0.2 | 2.3$\pm$0.6 | IFMR
Ross 640 | 1628+3646 | WD 1626+368 | 0.58$\pm$0.03 | Phot | 1 | 1.1$\pm$0.6 | 1.8$\pm$0.1 | 1.5$\pm$0.6 | IFMR
G 138-38 | 1632+0851 | WD 1630+089 | 0.59$\pm$0.15 | Spec | 2 | 1.3$\pm$1.9 | 1.9$\pm$0.1 | 1.6$\pm$1.9 | IFMR
EGGR 258 | 1634+5710 | WD 1633+572 | 0.57$\pm$0.04 | Phot | 1 | 1.0$\pm$0.7 | 1.7$\pm$0.1 | 1.4$\pm$0.7 | IFMR
PG 1633+434 | 1635+4317 | WD 1633+433 | 0.68$\pm$0.04 | Phot | 1 | 2.4$\pm$0.8 | 2.8$\pm$0.2 | 2.6$\pm$0.8 | IFMR
DN Dra | 1648+5903 | WD 1647+591 | 0.77$\pm$0.03 | Astero | 14 | 3.1$\pm$1.5 | 3.3$\pm$0.2 | 3.2$\pm$1.5 | IFMR
UCAC4 317-104829 | 1706$-$2643 | (WD 1703$-$267) | 0.808$\pm$0.009 | Gaia H-atm | 8 | 3.3$\pm$1.5 | 3.5$\pm$0.2 | 3.4$\pm$1.5 | IFMR
EGGR 494 | 1708+0257 | WD 1705+030 | 0.68$\pm$0.09 | Phot | 1 | 2.4$\pm$1.3 | 2.8$\pm$0.2 | 2.6$\pm$1.3 | IFMR
G 203-47 B | 1709+4340 | WD 1708+437 |  $>$0.50 | Dyn+mod | 22 |  $>$1.4 |  $>$2.0 |  $>$1.4 | IFMR
PM J17476$-$5436 | 1747$-$5436 |  $\cdots$ | 0.48$\pm$0.01 | Gaia He-atm | 8 |  $\cdots$ |  $\cdots$ | 1.0$\pm$0.1 | Low
EGGR 372 | 1748+7052 | WD 1748+708 | 0.79$\pm$0.01 | Phot | 1 | 3.2$\pm$1.5 | 3.4$\pm$0.2 | 3.3$\pm$1.5 | IFMR
EGGR 199 | 1749+8246 | WD 1756+827 | 0.55$\pm$0.13 | Spec | 7 |  $\cdots$ | 1.5$\pm$0.1 | 1.5$\pm$0.1 | IFMR
LSR J1817+1328 | 1817+1328 | WD 1814+134 | 0.68$\pm$0.02 | Phot | 1 | 2.4$\pm$0.7 | 2.8$\pm$0.2 | 2.6$\pm$0.7 | IFMR
G 227-28 | 1821+6101 | WD 1820+609 | 0.56$\pm$0.05 | Phot | 1 | 0.9$\pm$0.8 | 1.6$\pm$0.1 | 1.2$\pm$0.8 | IFMR
EGGR 176 | 1824$-$1308 | WD 1821$-$131 | 1.06$\pm$0.07 | Spec | 3 | 5.5$\pm$1.3 | 6.1$\pm$0.4 | 5.8$\pm$1.3 | IFMR
UCAC4 508-079937 | 1825+1135 |  $\cdots$ | 0.51$\pm$0.01 | Gaia H-atm | 8 |  $\cdots$ | 1.1$\pm$0.1 | 1.1$\pm$0.1 | IFMR
EGGR 374 | 1830+5447 | WD 1829+547 | 0.90$\pm$0.07 | Phot | 1 | 4.0$\pm$1.1 | 5.0$\pm$0.3 | 4.5$\pm$1.2 | IFMR
LAWD 73 | 1900+7039 | WD 1900+705 | 0.93$\pm$0.02 | Phot | 1 | 4.3$\pm$1.0 | 5.3$\pm$0.3 | 4.8$\pm$1.0 | IFMR
EGGR 375 | 1918+3843 | WD 1917+386 | 0.75$\pm$0.04 | Phot | 1 | 3.0$\pm$1.5 | 3.2$\pm$0.2 | 3.1$\pm$1.5 | IFMR
LAWD 74 | 1920$-$0740 | WD 1917$-$077 | 0.62$\pm$0.02 | Phot | 1 | 1.6$\pm$0.6 | 2.2$\pm$0.2 | 1.9$\pm$0.6 | IFMR
UCAC4 482-095741 | 1921+0613 |  $\cdots$ | 0.68$\pm$0.01 | Gaia H-atm | 8 | 2.4$\pm$0.6 | 2.8$\pm$0.2 | 2.6$\pm$0.7 | IFMR
GD 219 | 1921+1440 | WD 1919+145 | 0.74$\pm$0.03 | Spec | 1 | 3.0$\pm$1.4 | 3.1$\pm$0.2 | 3.1$\pm$1.5 | IFMR
PY Vul | 1937+2743 | WD 1935+276 | 0.66$\pm$0.02 | Astero | 15 | 2.1$\pm$0.6 | 2.7$\pm$0.2 | 2.4$\pm$0.7 | IFMR
LAWD 79 | 1956-0102 | WD 1953$-$011 | 0.79$\pm$0.13 | Spec | 7 | 3.2$\pm$1.7 | 3.4$\pm$0.2 | 3.3$\pm$1.7 | IFMR
Wolf 1130 B | 2005+5426 | WD 2003+542 | 1.24$\pm$0.17 | Orbit | 16 | 7.2$\pm$2.0 | 7.2$\pm$0.4 | 7.2$\pm$2.1 | IFMR
EGGR 498 | 2005$-$1056 | WD 2002$-$110 | 0.72$\pm$0.01 | Phot | 1 | 2.9$\pm$1.4 | 3.0$\pm$0.2 | 3.0$\pm$1.4 | IFMR
CD$-$30 17706 | 2010$-$3013 | WD 2007$-$303 | 0.60$\pm$0.02 | Spec | 1 | 1.4$\pm$0.5 | 2.0$\pm$0.2 | 1.7$\pm$0.6 | IFMR
SCR J2012$-$5956 | 2012$-$5956 | WD 2008$-$600 | 0.44$\pm$0.01 | Phot | 1 |  $\cdots$ |  $\cdots$ | 1.9$\pm$0.9 | Ultra-low
EC 20173$-$3036 | 2020$-$3027 |  $\cdots$ | 0.75$\pm$0.01 | Gaia H-atm* | 8 | 3.0$\pm$1.5 | 3.2$\pm$0.2 | 3.1$\pm$1.5 | IFMR
HD 340611 | 2034+2503 | WD 2032+248 | 0.64$\pm$0.03 | Spec | 1 | 1.9$\pm$0.7 | 2.4$\pm$0.2 | 2.2$\pm$0.7 | IFMR
EGGR 140 | 2044$-$6805 | WD 2039$-$682 | 0.98$\pm$0.03 | Spec | 1 | 4.8$\pm$1.0 | 5.6$\pm$0.4 | 5.2$\pm$1.1 | IFMR
EGGR 261 | 2049+3728 | WD 2047+372 | 0.81$\pm$0.03 | Spec | 1 | 3.3$\pm$1.5 | 3.5$\pm$0.2 | 3.4$\pm$1.5 | IFMR
G 187-8 A | 2050+2630 | WD 2048+263 | 0.24$\pm$0.04eeQuoted mass estimate assumes the system is a single white dwarf. | Phot | 1 |  $\cdots$ |  $\cdots$ | 1.9$\pm$0.9 | Ultra-low
G 187-8 B | 2050+2630 |  $\cdots$ |  $\cdots$ |  $\cdots$ | 1 |  $\cdots$ |  $\cdots$ | 1.9$\pm$0.9 | Ultra-low
UCAC4 325-215293 | 2052$-$2504 | (WD 2049$-$253) | 0.47$\pm$0.01 | Gaia He-atm | 8 |  $\cdots$ |  $\cdots$ | 1.0$\pm$0.1 | Low
Ross 193 B | 2056$-$0450 | WD 2054$-$050 | 0.49$\pm$0.01 | Spec | 23 |  $\cdots$ |  $\cdots$ | 1.0$\pm$0.1 | Low
WT 765 | 2101$-$4906 |  $\cdots$ | 0.53$\pm$0.01 | Gaia H-atm | 8 |  $\cdots$ | 1.3$\pm$0.1 | 1.3$\pm$0.1 | IFMR
LAWD 83 | 2113$-$8149 | WD 2105$-$820 | 0.78$\pm$0.10 | Spec | 7 | 3.2$\pm$1.6 | 3.4$\pm$0.2 | 3.3$\pm$1.6 | IFMR
EGGR 378 | 2118+5412 | WD 2117+539 | 0.56$\pm$0.03 | Spec | 1 | 0.9$\pm$0.6 | 1.6$\pm$0.1 | 1.2$\pm$0.6 | IFMR
$\nu$ Oct B | 2141$-$7723 |  $\cdots$ | 0.55$\pm$0.05 | Dyn+mod | 20 |  $\cdots$ | 1.5$\pm$0.1 | 1.5$\pm$0.1 | IFMR
L 570-26 | 2141$-$3300 | WD 2138$-$332 | 0.70$\pm$0.02 | Phot | 1 | 2.6$\pm$0.7 | 2.9$\pm$0.2 | 2.8$\pm$0.7 | IFMR
EGGR 148 | 2142+2059 | WD 2140+207 | 0.48$\pm$0.04 | Phot | 1 |  $\cdots$ |  $\cdots$ | 1.0$\pm$0.1 | Low
PHL 1716 | 2143$-$0659 |  $\cdots$ | 0.87$\pm$0.01 | Gaia H-atm | 8 | 3.7$\pm$0.9 | 4.5$\pm$0.3 | 4.1$\pm$1.0 | IFMR
UCAC4 747-070768 | 2151+5917 | (WD 2150+591) | 0.57$\pm$0.01 | Gaia H-atm | 8 | 1.0$\pm$0.4 | 1.7$\pm$0.1 | 1.4$\pm$0.5 | IFMR
WG 39 | 2157$-$5100 | WD 2154$-$512 | 0.60$\pm$0.04 | Phot | 1 | 1.4$\pm$0.7 | 2.0$\pm$0.2 | 1.7$\pm$0.7 | IFMR
CD Oct | 2204$-$7513 | WD 2159$-$754 | 0.92$\pm$0.04 | Spec | 1 | 4.2$\pm$1.0 | 5.3$\pm$0.3 | 4.7$\pm$1.1 | IFMR
WD 2211$-$392 | 2214$-$3859 | WD 2211$-$392 | 0.80$\pm$0.04 | Phot | 1 | 3.3$\pm$1.5 | 3.5$\pm$0.2 | 3.4$\pm$1.5 | IFMR
SCR J2230$-$7515 | 2230$-$7515 | WD 2226$-$755 | 0.58$\pm$0.00 | Phot | 1 | 1.1$\pm$0.4 | 1.8$\pm$0.1 | 1.5$\pm$0.4 | IFMR
SCR J2230$-$7513 | 2230$-$7513 | WD 2226$-$754 | 0.58$\pm$0.00 | Phot | 1 | 1.1$\pm$0.4 | 1.8$\pm$0.1 | 1.5$\pm$0.4 | IFMR
EGGR 155 | 2249+2236 | WD 2246+223 | 1.11$\pm$0.21 | Spec | 7 | 6.0$\pm$2.3 | 6.4$\pm$0.4 | 6.2$\pm$2.3 | IFMR
EGGR 283 A | 2251+2939 | WD 2248+293 | 0.35$\pm$0.07eeQuoted mass estimate assumes the system is a single white dwarf. | Phot | 1 |  $\cdots$ |  $\cdots$ | 1.9$\pm$0.9 | Ultra-low
EGGR 283 B | 2251+2939 |  $\cdots$ |  $\cdots$ |  $\cdots$ | 1 |  $\cdots$ |  $\cdots$ | 1.9$\pm$0.9 | Ultra-low
EGGR 453 | 2253$-$0646 | (WD 2251$-$070) | 0.58$\pm$0.03 | Phot | 1 | 1.1$\pm$0.6 | 1.8$\pm$0.1 | 1.5$\pm$0.6 | IFMR
LSPM J2309+5506E | 2309+5506 | WD 2307+548 | 0.59$\pm$0.15 | Spec | 2 | 1.3$\pm$1.9 | 1.9$\pm$0.1 | 1.6$\pm$1.9 | IFMR
ZZ Psc | 2328+0514 | WD 2326+049 | 0.593$\pm$0.012 | Astero | 15 | 1.3$\pm$0.5 | 1.9$\pm$0.1 | 1.6$\pm$0.5 | IFMR
GD 1212 | 2338$-$0741 | WD 2336$-$079 | 0.62$\pm$0.03 | Astero | 17 | 1.6$\pm$0.6 | 2.2$\pm$0.2 | 1.9$\pm$0.7 | IFMR
LAWD 93 | 2343+3232 | WD 2341+322 | 0.65$\pm$0.11 | GravRed | 5 | 2.0$\pm$1.5 | 2.6$\pm$0.2 | 2.3$\pm$1.5 | IFMR

References. — References for final mass measurements and estimates: (1)
Giammichele et al. (2012), (2) Limoges et al. (2015), (3) Gianninas et al.
(2011), (4) Maxted et al. (2002), (5) Silvestri et al. (2001), (6) Brandt et
al. (2019), (7) Bédard et al. (2017), (8) Gentile Fusillo et al. (2021), (9)
Mason et al. (2017), (10) Sahu et al. (2017), (11) Joyce et al. (2018), (12)
Gentile Fusillo et al. (2019), (13) Liebert et al. (2013), (14) Romero et al.
(2013), (15) Romero et al. (2012), (16) Mace et al. (2018), (17) Hermes et al.
(2014), (18) Baroch et al. (2021), (19) Zhuchkov et al. (2012), (20) Ramm et
al. (2021), (21) Winters et al. (2020), (22) Delfosse et al. (1999a), (23)
Blouin et al. (2019).

Note. — Codes for mass methods: Direct measurements of final mass (in
boldface) – Accel = dynamical acceleration, Astero = asteroseismology, GravRed
= gravitational redshift, Lensing = gravitational lensing, Orbit = dynamical
orbital analysis. Estimates of final mass (in normal font) – Dyn+mag = partial
dynamical orbital analysis and absolute magnitude, Dyn+mod = partial dynamical
orbital analysis used along with modeling, Gaia H-atm = Gaia-centric hydrogen-
atmosphere solution, Gaia H-atm* = assumes a Gaia-centric hydrogen-atmosphere
solution even though no spectrum is available, Gaia He-atm = Gaia-centric
helium-atmosphere solution, Phot = uses fit to the spectral energy
distribution, Spec = uses fit to the spectrum. Method for estimating the
initial mass – IFMR = initial-to-final-mass relation, Low = arbitrary
assignment for low final masses $0.45\leq M<0.56M_{\odot}$, Ultra-low =
special assignment for ultra-low final masses $M<0.45M_{\odot}$ (see text for
details), Conjecture = case-by-case handling (see text for details).

There are a number of caveats that must be addressed before proceeding with
the initial mass determinations, however: (1) The parameterization of the
initial-to-final mass relation is heavily reliant on white dwarfs with
hydrogen lines in their spectra, but roughly half of the white dwarfs in the
20-pc census have no hydrogen lines and are classified as DB (helium lines),
DC (no lines), DQ (carbon bands), or DZ (metal lines). Are such objects
expected to follow the same initial-to-final mass relation as the DA white
dwarfs? (2) Several 20-pc white dwarfs have masses lower than the
$0.50M_{\odot}$ lower bound of current initial-to-final mass relations. How do
we address this complication? (3) Given the age of the Milky Way, a white
dwarf with mass below ${\sim}0.45M_{\odot}$ should not yet exist (Fontaine et
al. 2001, Sun & Arras 2018), as there has been insufficient time for its
single-star progenitor to have evolved off the main sequence. Nonetheless, a
few 20-pc white dwarfs have final masses so low that they imply cooling ages
longer than a Hubble time. How do we interpret this seeming contradiction? We
address these issues below.

For issue (1), consider that within the 20-pc census, the most massive white
dwarfs323232One possible exception to this is Wolf 1130 B, which is believed
to have a progenitor mass near $8M_{\odot}$ and an oxygen-neon core (Mace et
al. 2018). are akin to Sirius B and have initial mass estimates of $\sim$5
$M_{\odot}$. With this in mind, it is helpful to summarize some relevant facts
about white dwarfs, as stated in the excellent review by Fontaine et al.
(2001). The alpha process (carbon burning) occurs only in stars more massive
than $\sim$8 $M_{\odot}$333333This is close to the lower mass boundary of
neutron star formation, which is believed to be $\sim 8-11M_{\odot}$ (Woosley
et al. 2002), so all white dwarfs within 20 pc of the Sun will have stopped
their thermonuclear burning at the triple-alpha process (helium burning).
These white dwarfs will have a core made primarily of carbon and oxygen and
will be compositionally stratified, with a helium-rich envelope and a separate
hydrogen-rich envelope. These two layers are very thin, but their opacity
regulates the core’s energy output. As such, these thin layers play a critical
role in determining white dwarf cooling times.

Current observational evidence shows that white dwarfs evolve
spectroscopically as they cool, sometimes appearing as hydrogen-atmosphere
stars (DA) and sometimes as helium-atmosphere stars (DB and other
classes343434The other non-DA spectral classes of white dwarfs are also
thought to be helium-atmosphere stars but with different temperatures and/or
trace pollutants than normal DB stars, with DO stars (ionized helium present)
at the hot end of the evolutionary sequence and DC, DQ, and DZ stars at the
cool end (figure 4 of Fontaine et al. 2013).), meaning that some unknown
process – convection, diffusion, and/or mixing – reorders the two outer
envelopes over time. Specifically, the ratio of helium-atmosphere to hydrogen-
atmosphere white dwarfs changes as a function of effective temperature, with
relatively few helium-atmosphere stars being identified in the range
$30,000<T_{\rm eff}<45,000$K (Shibahashi 2007) and relatively few helium-
atmosphere stars again being found in the cooler range $5000<T_{\rm
eff}<6000$K (Bergeron et al. 1997). This implies that most nearby white dwarfs
evolve as hydrogen-atmosphere stars, an assertion that Fontaine et al. (2001)
say is further bolstered by the fact that the observational data and
theoretical predictions – particularly with regards to the white dwarf
luminosity function – are in excellent agreement. We are therefore confident
in using the initial-to-final mass relations (IFMR) to estimate the initial
masses of 20-pc objects having final masses $>0.50M_{\odot}$. These initial
mass estimates are listed in Table 6.1.1 and are noted with "IFMR" in the
final column.

For issue (2), we note that of the 161 white dwarfs in Table 6.1.1, only
twelve have $0.45\leq M<0.50M_{\odot}$. An object at the upper end of this
range would have a predicted initial mass of $<0.9M_{\odot}$ according to the
initial-to-final mass relation of Cummings et al. (2018) (their equation 4) or
$\sim 1.0{\pm}0.1M_{\odot}$ according to the relation of El-Badry et al.
(2018) (their table 1 and their figure 3). Additionally, the lowest mass
object to have evolved off the main sequence during the lifetime of the
Universe (13.8 Gyr; Planck Collaboration et al. 2020) is predicted353535Below
masses of $\sim 5M_{\odot}$, objects of sub-solar metallicity are predicted to
evolve more rapidly than those of solar metallicity (figure 5 of Mowlavi et
al. 1998; see also figure 11 of Mowlavi et al. 2012). However, very few stars
in the solar neighborhood have metallicities significantly below solar. See
Section 6.2.2. to have an initial mass of $10^{-log(1.38)/2.5}=0.9M_{\odot}$
(equation 1.88 of Hansen & Kawaler 1994). Thus, we can consider all twelve of
these objects to have initial masses of $\sim 1.0{\pm}0.1M_{\odot}$, and these
low-mass estimates are noted by "Low" in the final column of Table 6.1.1. Non-
evolved late-F and early-G dwarfs in this same mass range significantly
outnumber (by $\sim 5\times$) the lowest mass white dwarf progenitors, so this
part of the white dwarf population makes only a small contribution to this
mass slice of the initial mass function anyway.

For issue (3), a small number of white dwarfs in the census have quoted masses
below ${\sim}0.45M_{\odot}$. There are two scenarios that can explain the
existence of such objects. The first is that unresolved white dwarf doubles
will be misinterpreted as being overluminous because, at a fixed value of
$T_{\rm eff}$, the Stefan-Boltzmann law will imply a falsely large radius,
which results in an erroneously small mass. Of the ultra-low mass systems
within 20 pc, three (LP 991-16, G 187-8, and EGGR 283) were suspected to be
double degenerate systems by Giammichele et al. (2012), and those predictions
appear to have been verified by Gaia DR3: LP 991-16 and G 187-8 both fall
significantly above the white dwarf sequence on the Gaia-based color-magnitude
diagram, and EGGR 283 is shown to be an astrometric double with an orbital
period of 278.0 d. The masses quoted for these objects in Table 6.1.1 are
therefore biased low, and these white dwarfs likely have masses
$>0.45M_{\odot}$. A fourth system, LP 941-19, has a Gaia DR3 RUWE value of
1.817 and may yet prove to be a double star in its own right and thus also
have a biased mass estimate.

The second scenario is that common-envelope mergers or episodic mass loss can
produce remnants with pure helium cores (Serenelli et al. 2001), and for these
the cooling times are much longer than those of white dwarfs with carbon-
oxygen cores, particularly if the hydrogen envelope is massive enough for its
own sustained burning (Alberts et al. 1996, Sarna et al. 1999). This scenario
likely explains the low mass estimate for SCR J2012$-$5956, which has a Gaia
DR3 RUWE value of 0.98 and is likely to be a single star.

As Giammichele et al. (2012) explain, without knowing more about the companion
objects in the double degenerate systems, determining new mass estimates for
the individual objects is not straightforward. Whether or not the other ultra-
low-mass systems are the products of a double-object merger or mass loss from
a single star is also guesswork. For these seven white dwarf systems, we have
estimated initial masses as follows. We know only that their initial masses
likely fall between $\sim$0.9 and $\sim$8.0 $M_{\odot}$. Because the median
initial mass of 20-pc white dwarfs with solid mass estimates is
$1.9M_{\odot}$, we arbitrarily assign each a mass of $1.9{\pm}0.9M_{\odot}$
that, at 1-sigma, encompasses the initial mass range spanned by 72% of the
20-pc white dwarf sample. These estimates are marked with "Ultra-low" in the
last column of Table 6.1.1.

Finally, we note three objects in Table 6.1.1 whose white dwarf natures are
more speculative. Each of these objects is discussed individually below:

  * •

$\chi$ Eri B: Both Fuhrmann & Chini (2012a) and Fuhrmann et al. (2016) posit
that the appreciable X-ray luminosity coming from this system emanates not
from the G9 subgiant primary (Gray et al. 2006), $\chi$ Eri A, but from the
secondary, $\chi$ Eri B. Although this companion is seven magnitudes fainter
at $V$ band, they stipulate that it could account for the anomalous X-ray flux
if it were a white dwarf.

  * •

$\alpha$ For Bb: The primary in this system, $\alpha$ For A, is an F6 dwarf
(Gray et al. 2006) and the secondary, $\alpha$ For B, is a G7 dwarf (Corbally
1984). This B component was found to be a 3.75-d radial velocity binary by
Fuhrmann et al. (2016). No spectral lines are visible from the tertiary
component, meaning that its mass would have to be below $0.35M_{\odot}$ if it
were an M dwarf. However, Fuhrmann et al. (2016) find that the spectrum of
$\alpha$ For Ba is enhanced in barium content relative to other dwarfs of
similar spectral type and relative to its primary star, $\alpha$ For A. As
barium is an abundant product of nucleosynthesis on the red giant branch, the
authors speculate that $\alpha$ For Ba was polluted during the post-main
sequence evolution of $\alpha$ For Bb, now a white dwarf.

  * •

$\xi$ UMa Bb: This is part of a quintuple system. The primary, $\xi$ UMa A, is
an F8.5: dwarf (Keenan & McNeil 1989) and the secondary, $\xi$ UMa B, is a G2
dwarf (Keenan & McNeil 1989). Both components are spectroscopic binaries. The
A component is an RS CVn double (Samus’ et al. 2003), and the B component is a
well known double with a period of $\sim$4 d, as summarized in Fuhrmann
(2008). A distant, co-moving fifth member of this system, the T8.5 dwarf WISE
J111838.70+312537.9, has also been identified; Wright et al. 2013. As with
$\alpha$ For Ba discussed above, $\xi$ UMa Ba has an enhanced barium
abundance, leading Fuhrmann et al. (2016) to speculate that $\xi$ UMa Bb is a
white dwarf and the donor responsible for the extra barium content.

For these three white dwarfs, we also arbitrarily assign each an initial mass
of $1.9{\pm}0.9M_{\odot}$, as done for the objects with ultra-low mass
estimates. Estimates for these three white dwarfs are marked with "Conjecture"
in the last column of Table 6.1.1.

####  6.1.2 Giants and subgiants

There are a number of objects in the 20-parsec census that have evolved off
the main sequence but have not yet become white dwarfs. Table 6.1.2 includes
all objects in Table 3.1 that have a luminosity class more luminous than V
and/or fall in a locus on the absolute magnitude vs. color diagrams that
identifies them as post-main sequence stars.

Several of these have direct mass measurements either from orbital dynamics or
asteroseismology. The rest have had their masses estimated from other methods,
primarily via comparison of their placement on the HR diagram in relation to
modeled evolutionary tracks or via fits of their spectra to atmospheric
models. For some objects with IV-V or IV luminosity classes, other published
spectral types indicate a V luminosity class or their placement on the HR
diagram suggests a main sequence star. This is reflected in the mass estimates
given in Table 6.1.2.

Table 15: Giants and Subgiants in the 20-pc Census

Name | Abbrev. Coords | Sp. TypeaaReferences for Sp. Type can be found in Table 3.1. | MassbbMethods in bold involve direct mass measurements. | Mass | Method
---|---|---|---|---|---
| J2000 |  | ($M_{\odot}$) | Ref. |
(1) | (2) | (3) | (4) | (5) | (6)
$\beta$ Cas | 0009+5908 | F2 III | 1.91$\pm$0.02 | 9 | Interferometry + rapid-rotation models
$\alpha$ Tri A | 0153+2934 | F6 IV | 1.70 | 15 | Spectral fit to atmospheric models
$\chi$ Eri A | 0155$-$5136 | G9 IV | 1.58 | 13 | Placement on evolutionary tracks
10 Tau | 0336+0024 | F9 IV-V | 1.139$\pm$0.016 | 23 | Interferometry + model isochrones
$\delta$ Eri | 0343$-$0945 | K1 III-IV | 1.33$\pm$0.07 | 8 | Asteroseismology
$\tau^{6}$ Eri | 0346$-$2314 | F5 IV-V | 1.44$\pm$0.13 | 13 | Placement on evolutionary tracks
$\epsilon$ Ret A | 0416$-$5918 | K2 III | 1.48 | 25 | Spectral fit to atmospheric models
HD 283750 Aa | 0436+2707 | K3 IV ke | 0.84$\pm$0.19 | 15 | Placement on evolutionary tracks
$\pi^{3}$ Ori | 0449+0657 | F6 IV-V | 1.283$\pm$0.006 | 23 | Interferometry + model isochrones
$\alpha$ Aur Aa | 0516+4559 | G1 III | 2.569$\pm$0.007 | 27 | Orbital dynamics
$\alpha$ Aur Ab | 0516+4559 | K0 III | 2.483$\pm$0.007 | 27 | Orbital dynamics
$\xi$ Gem | 0645+1253 | F5 IV-V | 1.706$\pm$0.012 | 23 | Interferometry + model isochrones
HD 53143 | 0659$-$6120 | K0 IV-V (k) | 1.0 | 22 | Placement on color-magnitude diagram
$\alpha$ Gem Aa | 0734+3153 | A1.5 IV+ | 2.98 | 26 | Estimated from spectral type/color
$\beta$ Gem | 0745+2801 | K0 III | 1.91$\pm$0.09 | 16 | Asteroseismology
$\rho$ PupccLuminosity class II suggests a more evolved state for this star than its placement on the HR diagram – subgiant or giant – attests. The cause for this "anomalous luminosity effect" is unknown but is a feature of the $\rho$ Puppis class of pulsators (Gray & Corbally 2009). | 0807$-$2418 | F5II kF2II mF5II | 1.9$\pm$0.1 | 1 | Placement of general $\rho$ Pup class on evolutionary tracks
HD 73752 Aa | 0839$-$2239 | G5 IV | 1.21 | 14 | Spectral fit to atmospheric models
$\rho^{1}$ Cnc A | 0852+2819 | K0 IV-V | 0.91$\pm$0.02 | 28 | Placement on evolutionary tracks
10 UMa A | 0900+4146 | F5 IV-V | 1.396$\pm$0.002 | 4 | Orbital dynamics
HD 78366ddFuhrmann (2008) believes this object is young, not evolved. | 0908+3352 | G0 IV-V | 1.08 | 15 | Spectral fit to atmospheric models
$\tau^{1}$ Hya Aa | 0929$-$0246 | F5.5 IV-V | 1.20 | 11 | Placement on evolutionary tracks
$\theta$ UMa A | 0932+5140 | F5.5 IV-V | 1.506$\pm$0.095 | 23 | Interferometry + model isochrones
15 LMi | 0948+4601 | G0 IV-V | 1.11+0.15 | 7 | Spectral fit to atmospheric models
$\delta$ Leo | 1114+2031 | A5 IV(n) | 2.061$\pm$0.006 | 23 | Interferometry + model isochrones
$\beta$ Vir | 1150+0145 | F8.5 IV-V | 1.42$\pm$0.08 | 8 | Asteroseismology
HD 104304 A | 1200$-$1026 | G8 IV | 0.98 | 25 | Spectral fit to atmospheric models
e Vir Aa | 1316+0925 | G0 IV | 1.22 | 12 | Placement on evolutionary tracks
70 Vir | 1328+1346 | G5 VffSome references classify this object as G4 V-IV (e.g., Strassmeier et al. 2018), and its placement just above the main sequence indicates that it may just be moving into a later stage of evolution. | 1.14$\pm$0.08 | 31 |  $L_{bol}$ \+ $T_{\rm eff}$ \+ spectrum-based log(g)
$\tau$ Boo A | 1347+1727 | F7 IV-V | 1.34 | 25 | Spectral fit to atmospheric models
$\eta$ Boo Aa | 1354+1823 | G0 IV | 1.77$\pm$0.11 | 8 | Asteroseismology
$\theta$ Cen | 1406$-$3622 | K0 III | 1.27 | 13 | Placement on evolutionary tracks
$\alpha$ Boo | 1415+1910 | K0 III | 0.80$\pm$0.20ggEvolutionary models suggest a mass of $1.08{\pm}0.06M_{\odot}$, toward the upper end of the range deduced from asteroseismology, and a relatively old age of $7.1^{+1.5}_{-1.2}$ Gyr (Ramírez & Allende Prieto 2011). (An object with a mass of $\sim$0.90 $M_{\odot}$ has a main sequence lifetime exceeding a Hubble time and will not yet have evolved to a giant state; see section 1.7 of Hansen & Kawaler 1994.) | 19 | Asteroseismology
HD 125072 | 1419$-$5922 | K3 IV | 0.88 | 25 | Spectral fit to atmospheric models
$\alpha$ Cen B | 1439$-$6050 | K2 IV C2+1** | 0.909$\pm$0.003 | 30 | Orbital dynamics
HD 130948 A | 1450+2354 | F9 IV-V | 1.18$\pm$0.16 | 7 | Spectral fit to atmospheric models
$\lambda$ Ser | 1546+0721 | G0 IV-V | 1.15$\pm$0.15 | 7 | Spectral fit to atmospheric models
HD 140901 A | 1547$-$3754 | G7 IV-V | 0.99 | 25 | Spectral fit to atmospheric models
14 Her | 1610+4349 | K0 IV-V | 0.73$\pm$0.10 | 7 | Spectral fit to atmospheric models
$\zeta$ Her A | 1641+3136 | G2 IV | 1.04$\pm$0.03 | 20 | Orbital dynamics
$\epsilon$ Sco A | 1650$-$3417 | K1 III | 1.4$\pm$0.1 | 18 | Asteroseismology
HD 154088 | 1704$-$2834 | K0 IV-V | 0.92 | 25 | Spectral fit to atmospheric models
HD 158614 A | 1730$-$0103 | G9- IV-V Hdel1 | 0.963$\pm$0.005 | 4 | Orbital dynamics
$\alpha$ Oph A | 1734+1233 | A5 IV nn | 2.20$\pm$0.06 | 17 | Orbital dynamics
26 Dra A | 1734+6152 | G0 IV-V | 1.06 | 10 | Orbital dynamics + astrophysical assumptions
$\mu$ Ara | 1744$-$5150 | G3 IV-V | 1.21$\pm$0.13 | 8 | Asteroseismology
$\mu^{1}$ Her A | 1746+2743 | G5 IV | 1.10${}^{+0.11}_{-0.06}$ | 21 | Asteroseismology
$\eta$ Ser | 1821$-$0253 | K0 III-IV | 1.45$\pm$0.21 | 8 | Asteroseismology
110 Her | 1845+2032 | F5.5 IV-V | 1.422$\pm$0.009 | 23 | Interferometry + model isochrones
b Aql | 1924+1156 | G7 IV Hdel1 | 1.186$\pm$0.015 | 23 | Interferometry + model isochrones
$\delta$ Aql Aa | 1925+0306 | F1 IV-V(n) | 1.45 | 11 | Placement on evolutionary tracks
HD 188088 Aa | 1954$-$2356 | K2 IV (k) | 0.85 | 29 | Orbital dynamics + other assumptionseeFekel et al. (2017) calculates minimum masses ($M\sin^{3}{i}$) from orbital dynamics of 0.8463$\pm$0.0014 $M_{\odot}$ for HD 188088 A and 0.8316$\pm$0.014 $M_{\odot}$ for HD 188088 B. These authors believe that both components are normal K dwarfs, the subgiant classification likely resulting from the slight metal richness of these stars. Given that these minimum masses are close to the mass expected for this dwarf class, the inclination is suspected of being near $90^{\circ}$ despite the lack of eclipses in the system.
$\beta$ Aql A | 1955+0624 | G9.5 IV | 1.26$\pm$0.18 | 8 | Asteroseismology
HD 190360 | 2003+2953 | G7 IV-V | 0.92$\pm$0.12 | 7 | Spectral fit to atmospheric models
$\delta$ Pav | 2008$-$6610 | G8 IV | 1.07$\pm$0.13 | 8 | Asteroseismology
$\eta$ Cep | 2045+6150 | K0 IV | 1.6 | 2 | Placement on evolutionary tracks
$\nu$ Oct A | 2141$-$7723 | K1 III | 1.6$\pm$0.1 | 24 | Placement on evolutionary tracks
$\delta$ Cap Aa | 2147$-$1607 | kA5hF0mF2 III | 2.0 | 6 | Orbital dynamics
$\gamma$ Cep A | 2339+7737 | K1 III | 1.294$\pm$0.081 | 5 | Orbital dynamics

References. — References for the mass measurements and estimates: (1) Abt
2017, (2) Affer et al. 2005, (3) Allende Prieto & Lambert 1999, (4) Andrade
2019, (5) Mugrauer et al. 2022, (6) Batten & Fletcher 1992, (7) Brewer et al.
2016, (8) Bruntt et al. 2010. (9) Che et al. 2011, (10) Cvetkovic & Ninkovic
2010, (11) David & Hillenbrand 2015, (12) D’Orazi et al. 2017, (13) Fuhrmann &
Chini 2012a, (14) Fuhrmann et al. 2011a, (15) Fuhrmann 2008, (16) Hatzes et
al. 2012, (17) Gardner et al. 2021, (18) Kallinger et al. 2019, (19) Kallinger
et al. 2010, (20) Katoh et al. 2013, (21) Li et al. 2019, (22) Nielsen et al.
2019, (23) Boyajian et al. 2012, (24) Ramm et al. 2021, (25) Takeda et al.
2007, (26) Tokovinin 2008, (27) Torres et al. 2015, (28) von Braun et al.
2011, (29) Fekel et al. 2017, (30) Akeson et al. 2021, (31) Stassun et al.
2017.

One curious observation from the absolute magnitude vs. spectral type diagrams
of Figure 4 is the vertical locus of evolved stars lying well above the main
sequence but concentrated almost exclusively at a spectral class K0. On the
$M_{W4}$ vs. spectral type plot, for example, we find fifteen objects (not
counting the young, main sequence star $\beta$ Pic, whose unusual position is
caused by its debris disk) that lie more than one magnitude above the main
sequence for their spectral classes. These can be divided into a group of five
objects ($\beta$ Cas, $\rho$ Pup, $\alpha$ Tri, $\eta$ Boo, and $\zeta$ Her)
with classes between F2 and G2 and a group of ten objects ($\delta$ Eri,
$\beta$ Aql A, $\eta$ Cep, $\gamma$ Cep A, $\nu$ Oct A, $\eta$ Ser, $\beta$
Gem, $\theta$ Cen, $\epsilon$ Sco A, and $\alpha$ Boo) with a very narrow
range of types from G9.5 to K1. As we can see from Table 6.1.2, this first
group of five objects has a mass range of 1.0-1.9 $M_{\odot}$ and falls in a
locus in Figure 4 indicating evolution off the main sequence and onto the
evolutionary subgiant branch363636Note that the assigning of an object’s
subgiant or giant luminosity class via spectroscopic gravity diagnostics does
not necessarily equate to its presumed evolutionary status as a subgiant or
red giant branch star via its placement in color magnitude diagrams.. The
other group, of ten objects, is comprised of stars with an identical mass
range (1.0-1.9 $M_{\odot}$) that are now ascending the red giant branch.
Evolution along the subgiant branch is more rapid than the climb up the red
giant branch, explaining the overabundance of red giants (Tables 3-4 of Iben
1967). The fact that these latter stars are concentrated so narrowly in
spectral type is a consequence of the fact that at typical disk ages for
masses in this same range, the red giant branch is confined narrowly to a
temperature of $\sim$5000K (Figure 13 of Iben 1967, Figure 8 of Bressan et al.
2012), which is the temperature that corresponds to early-K giant spectral
classes (Figure 2 of Dyck et al. 1996, Figure 4 of Richichi et al. 1999, Table
10 of Heiter et al. 2015). Objects above this mass range ($>1.9M_{\odot}$) are
few in number in the 20-pc census ($<$3% of the total; see Table 3.1) and
evolve through their giant phases in no more than a few tens of Myr (Table 3
of Iben 1967); as a consequence of their rarity and rapid evolution, no such
earlier type giants are seen.

####  6.1.3 Brown dwarfs

Brown dwarfs follow no mass-luminosity relation because they constantly cool
over time. If the age of the brown dwarf is known, this can be used to
estimate the mass from evolutionary models, but age is a difficult parameter
to measure for non-youthful disk objects. We therefore must resort to
simulations to tease out information regarding the mass function. In
Kirkpatrick et al. (2019, 2021a), we took the empirical distribution of brown
dwarf effective temperatures and compared that to various predicted
temperature distributions modeled by taking the shape of the brown dwarf mass
function, the value of its low-mass cutoff, and the underlying evolutionary
model suite as free parameters. For the analysis of this paper, we will employ
those same methods, using an updated suite of predictions by Raghu et al.
(submitted).

Here, we compare the Kirkpatrick et al. (2021a) accounting of all 525 known
20-pc L, T, and Y dwarfs to that given in Table 3.1. Additions and
subtractions to this tally are listed in Table 16. We find that eight objects
have fallen out of the 20-pc sample, all because of new parallax measurements
or revised distance estimates that place them outside of 20 pc. On the other
hand, sixty-five objects are newly added. These additions include thirty-eight
new discoveries (thirty-seven by the Backyard Worlds citizen science group,
four of which are new companions), nine new companions recently announced in
the literature, one new companion announced here but found in Gaia, three new
published parallaxes with d$<$20 pc, twelve previously overlooked companions,
and two previously overlooked objects (DENIS J065219.7-253450, presumably due
to a transcription error, and SSSPM J1444-2019, whose subdwarf type had
earlier been updated from late-M to early-L). To facilitate analysis on the
revised $T_{\rm eff}$ distribution, we have listed in Table 16 the estimated
temperatures of each of the additions and subtractions. Further analysis can
be found in Section 7.

Table 16: Additions to and Subtractions from the 20-pc L, T, and Y Census of Kirkpatrick et al. (2021a) Object | Reason for Change | Reference |  $T_{\rm eff}$ (K)aaFor the new additions, temperature values with uncertainties of $\pm$88K were determined via the $M_{H}$ vs. $T_{\rm eff}$ relation in table 13 of Kirkpatrick et al. (2021a) and those with uncertainties of $\pm$134K or $\pm$79K were determined with the (assumed) spectral type vs. $T_{\rm eff}$ relations in the same table, unless otherwise noted. For subtractions, the temperature values are taken from table 11 of Kirkpatrick et al. (2021a).
---|---|---|---
(1) | (2) | (3) | (4)
Additions
WISE J003110.04+574936.3 B | New companion | Best et al. (2021) | 1275$\pm$200bbNo information is yet available on this companion, so the temperature estimate is set to cover the spectral type range from late-L to mid-T.
CWISE J003507.81$-$153233.5 | New discovery | This paper (Table 2) | 686$\pm$79
2MASSW J0036159+182110 B | Overlooked companion | Bernat et al. (2010) | 1125$\pm$79
CWISE J013343.58+803153.1 | New discovery | This paper (Table 2) | 1181$\pm$79
CWISE J014433.03$-$545545.5 | New discovery | This paper (Table 2) | 751$\pm$79
CWISER J021550.96+674017.2 | New companion to HD 13579 | This paper (Table 2 and Section 2.2) | 1125$\pm$79
CWISER J021612.11+423015.9 | New discovery | This paper (Table 2) | 460$\pm$79
L 440-30 B (0219$-$3646) | Overlooked companion | Kürster et al. (2008) | 1613$\pm$134
Ross 19 B (0219+3518) | New companion | Schneider et al. (2021) | 460$\pm$79
CWISE J032600.46+421058.5 | New discovery | This paper (Table 2) | 624$\pm$79
WISE J033605.05$-$014350.4 B | New companion | Calissendorff et al. (2023) | 325$\pm$79ccTemperature estimate is taken from Calissendorff et al. (2023), although the uncertainty has been inflated to match the typical uncertainties in Kirkpatrick et al. (2021a).
Wolf 227B (0352+1701) | Overlooked companion | Winters et al. (2018) | 2000$\pm$81hhWinters et al. (2018) conclude this is likely an early-L dwarf, so we set the temperature to cover the typical range for L0 to L4 dwarfs.
CWISE J035856.18+480244.9 | New discovery | This paper (Table 2) | 511$\pm$79
L 375-2 B (0432$-$3947) | New discovery | Silverstein et al. (2022) | 1200$\pm$333iiWe assume a huge temperature range to encompass the full substellar regime, as Silverstein et al. (2022) is able to provide only limited constraints on this companion.
LP 775-31 B (0435$-$1606) | Overlooked companion | Cortés-Contreras et al. (2017) | 1600$\pm$81jjThe $\Delta{I}$ magnitude from Cortés-Contreras et al. (2017) suggests a mid-L dwarf, so we estimate a temperature corresponding to this range.
Wolf 230 C (0507+1758) | Overlooked companion | Winters et al. (2020) | 1200$\pm$333kkWe assume a huge temperature range to encompass the full substellar regime, as Winters et al. (2020) is able to provide only an minimum mass of $\sim 44M_{Jup}$ for this companion.
CWISE J053046.20+440849.2 | New discovery | This paper (Table 2) | 1125$\pm$79
CWISE J060938.91+062513.2 | New discovery | This paper (Table 2) | 1420$\pm$134
DENIS J063001.4$-$184014 B | New companion | Sahlmann et al. (2021) | 1420$\pm$134
DENIS J063001.4$-$184014 (C) | New companion | Sahlmann et al. (2021) | 1420$\pm$134
WISEA J064750.85$-$154616.4 B | New companion | Best et al. (2021) | 1275$\pm$200bbNo information is yet available on this companion, so the temperature estimate is set to cover the spectral type range from late-L to mid-T.
DENIS J065219.7$-$253450 | Overlooked object | Gaia DR3 | 2196$\pm$88
PSO J103.0927+41.4601 B (0652+4127) | New companion | Best et al. (2021) | 1190$\pm$100ddNo information is yet available on this companion, so the temperature estimate is set to cover the spectral type range from early-T to mid-T.
CWISE J075227.38+053802.6 | New discovery | This paper (Table 2) | 1273$\pm$79