# 89 New Ultracool Dwarf Co-Moving Companions Identified With The Backyard
Worlds: Planet 9 Citizen Science Project

Austin Rothermich Department of Astrophysics, American Museum of Natural
History, Central Park West at 79th Street, NY 10024, USA  Department of
Physics, Graduate Center, City University of New York, 365 5th Ave., New York,
NY 10016, USA  Department of Physics and Astronomy, Hunter College, City
University of New York, 695 Park Avenue, New York, NY, 10065, USA  Backyard
Worlds: Planet 9  Jacqueline K. Faherty Department of Astrophysics, American
Museum of Natural History, Central Park West at 79th Street, NY 10024, USA
Backyard Worlds: Planet 9  Daniella Bardalez-Gagliuffi Department of
Astrophysics, American Museum of Natural History, Central Park West at 79th
Street, NY 10024, USA  Backyard Worlds: Planet 9  Adam C. Schneider US Naval
Observatory, Flagstaff Station, 10391 West Naval Observatory Road, Flagstaff,
AZ 86005, USA  Backyard Worlds: Planet 9  J. Davy Kirkpatrick IPAC, Mail Code
100-22, Caltech, 1200 E. California Blvd., Pasadena, CA 91125, USA  Backyard
Worlds: Planet 9  Aaron M. Meisner NSF’s National Optical-Infrared Astronomy
Research Laboratory, 950 N. Cherry Ave., Tucson, AZ 85719, USA  Backyard
Worlds: Planet 9  Adam J. Burgasser Department of Astronomy & Astrophysics,
University of California, San Diego, 9500 Gilman Drive, La Jolla, CA 92093,
USA  Backyard Worlds: Planet 9  Marc Kuchner Exoplanets and Stellar
Astrophysics Laboratory, NASA Goddard Space Flight Center, 8800 Greenbelt
Road, Greenbelt, MD 20771, USA  Backyard Worlds: Planet 9  Katelyn Allers
Physics and Astronomy Department, Bucknell University, 701 Moore Ave,
Lewisburg, PA 17837, USA  Jonathan Gagné Planètarium Rio Tinto Alcan, 4801
Pierre-de Coubertin Ave, Montreal, Quebec H1V 3V4, Canada  Backyard Worlds:
Planet 9  Dan Caselden Department of Astrophysics, American Museum of Natural
History, Central Park West at 79th Street, NY 10024, USA  Backyard Worlds:
Planet 9  Emily Calamari Department of Astrophysics, American Museum of
Natural History, Central Park West at 79th Street, NY 10024, USA  Department
of Physics, Graduate Center, City University of New York, 365 5th Ave., New
York, NY 10016, USA  Mark Popinchalk Department of Astrophysics, American
Museum of Natural History, Central Park West at 79th Street, NY 10024, USA
Department of Physics, Graduate Center, City University of New York, 365 5th
Ave., New York, NY 10016, USA  Department of Physics and Astronomy, Hunter
College, City University of New York, 695 Park Avenue, New York, NY, 10065,
USA  Genaro Suárez Department of Astrophysics, American Museum of Natural
History, Central Park West at 79th Street, NY 10024, USA  Roman Gerasimov
Center for Astrophysics and Space Science, University of California, San
Diego, 9500 Gilman Drive, La Jolla, CA 92093, USA  Department of Physics &
Astronomy, University of Notre Dame, Notre Dame, IN 46556, USA  Christian
Aganze Center for Astrophysics and Space Sciences, University of California,
San Diego, 9500 Gilman Drive, Mail Code 0424, La Jolla, CA 92093, USA
Department of Physics, Stanford University, Stanford CA 94305, USA  Emma
Softich Department of Astronomy & Astrophysics, University of California, San
Diego, 9500 Gilman Drive, La Jolla, CA 92093, USA  Chin-Chun Hsu Center for
Interdisciplinary Exploration and Research in Astrophysics (CIERA),
Northwestern University, 1800 Sherman, Evanston, IL, 60201, USA  Preethi
Karpoor  Department of Astronomy & Astrophysics, University of California, San
Diego, 9500 Gilman Drive, La Jolla, CA 92093, USA  Christopher A. Theissen
Department of Astronomy & Astrophysics, University of California, San Diego,
9500 Gilman Drive, La Jolla, CA 92093, USA  Jon Rees UCO/Lick Observatory,
7281 Mount Hamilton Road, Mount Hamilton, CA 95140, USA  Rosario Cecilio-
Flores-Elie  Department of Physics, Graduate Center, City University of New
York, 365 5th Ave., New York, NY 10016, USA  Michael C. Cushing  Ritter
Astrophysical Research Center, Department of Physics and Astronomy, University
of Toledo, 2801 W. Bancroft Street, Toledo, OH 43606, USA  Federico Marocco
IPAC, Mail Code 100-22, Caltech, 1200 E. California Blvd., Pasadena, CA 91125,
USA  Backyard Worlds: Planet 9  Sarah Casewell School of Physics and
Astronomy, University of Leicester, University Road, Leicester, LE1 7RH, UK
Backyard Worlds: Planet 9  Thomas P. Bickle School of Physical Sciences, The
Open University, Milton Keynes, MK7 6AA, UK  Backyard Worlds: Planet 9  Les
Hamlet  Backyard Worlds: Planet 9  Michaela B. Allen  NASA Goddard Space
Flight Center, Exoplanets and Stellar Astrophysics Laboratory, Code 667,
Greenbelt, MD 20771, USA  Paul Beaulieu  Backyard Worlds: Planet 9  Guillaume
Colin Backyard Worlds: Planet 9  Jean Marc Gantier Backyard Worlds: Planet 9
Leopold Gramaize Backyard Worlds: Planet 9  Peter Jalowiczor Backyard Worlds:
Planet 9  Martin Kabatnik Backyard Worlds: Planet 9  Frank Kiwy Backyard
Worlds: Planet 9  David W. Martin  Backyard Worlds: Planet 9  Billy Pendrill
Backyard Worlds: Planet 9  Ben Pumphrey Backyard Worlds: Planet 9  Arttu
Sainio Backyard Worlds: Planet 9  Jorg Schumann Backyard Worlds: Planet 9
Nikolaj Stevnbak Backyard Worlds: Planet 9  Guoyou Sun Backyard Worlds: Planet
9  Christopher Tanner  Backyard Worlds: Planet 9  Vinod Thakur  Backyard
Worlds: Planet 9  Melina Thevenot Backyard Worlds: Planet 9  Zbigniew Wedracki
Backyard Worlds: Planet 9

###### Abstract

We report the identification of 89 new systems containing ultracool dwarf
companions to main sequence stars and white dwarfs, using the citizen science
project Backyard Worlds: Planet 9 and cross-reference between Gaia and
CatWISE2020. Thirty-two of these companions and thirty-three host stars were
followed up with spectroscopic observations, with companion spectral types
ranging from M7-T9 and host spectral types ranging from G2-M9. These systems
exhibit diverse characteristics, from young to old ages, blue to very red
spectral morphologies, potential membership to known young moving groups, and
evidence of spectral binarity in 9 companions. Twenty of the host stars in our
sample show evidence for higher order multiplicity, with an additional 11 host
stars being resolved binaries themselves. We compare this sample’s
characteristics with those of the known stellar binary and exoplanet
populations, and find our sample begins to fill in the gap between directly
imaged exoplanets and stellary binaries on mass ratio-binding energy plots.
With this study, we increase the population of ultracool dwarf companions to
FGK stars by $\sim$42%, and more than triple the known population of ultracool
dwarf companions with separations larger than 1,000 au, providing excellent
targets for future atmospheric retrievals.

Brown Dwarfs, L Dwarfs, T Dwarfs, Astrometry, Spectroscopy, Stellar Rotation,
Wide Binary Stars

##  1 Introduction

Ultracool dwarfs, which consist of objects of effective temperatures ($T_{\rm
eff}$) below $\sim$2,700 K, include late-M spectral types and sub-stellar mass
objects known as brown dwarfs. Brown dwarfs bridge the mass regimes of stars
and giant exoplanets, with masses below the hydrogen burning limit ($\sim$75
$M_{\rm Jup}$) and above the deuterium burning limit ($\sim$13 $M_{\rm Jup}$)
(Kumar, 1963; Hayashi & Nakano, 1963; Burrows et al., 1997). Brown dwarfs
likely form through direct collapse and fragmentation of their natal molecular
cloud (Bate et al., 2002, 2003; Bate & Bonnell, 2005; Bate, 2009, 2012) and
therefore represent a low mass extension to star formation. The growing number
of isolated planetary mass sources discovered (for example, Gagné et al. 2015)
suggests this process is able to produce objects that overlap with the mass
regime of giant exoplanets. Many brown dwarf atmospheres also have
temperatures comparable to those found in exoplanets ($\sim$250–3000 K;
Kirkpatrick 2005), meaning both populations share atmospheric chemistry. While
the direct study of exoplanets is difficult, due in part to the contamination
from their host star’s light, many brown dwarfs are found without a
contaminating bright host, providing easier targets for spectroscopic
observations. Brown dwarfs are therefore ideal exoplanet analogs for the
testing and constraining of atmospheric models (Faherty et al., 2016).

Since the first confirmed discoveries (Rebolo et al., 1995; Nakajima et al.,
1995; Oppenheimer et al., 1995), numerous brown dwarfs have been identified
(see Kirkpatrick 2005). Characterizing brown dwarfs remains difficult though,
due to a degeneracy between mass, temperature, and age (Burrows et al., 2001).
The large number of different molecular species which form in brown dwarf
atmospheres presents another challenge, as it can be difficult to disentangle
the physical properties that dictate atmospheric chemistry and thereby sculpt
their highly structured spectra (Kirkpatrick, 2005; Helling & Casewell, 2014).
These factors make it a challenge to fully characterize an isolated brown
dwarf.

Brown dwarfs as co-moving companions to more easily characterized stars adopt
many of the properties of their hosts. Assuming the system to be coeval and
having formed from the same molecular cloud, properties like age and
metallicity that are measurable in the primary are expected to apply to the
companion as well. When in wide orbits (e.g. $>$100 AU), brown dwarf
companions within a few tens of parsecs can be fully resolved, allowing them
to be observed without interference from their host star, making these
valuable “benchmark” systems (e.g., Pinfield et al., 2006; Deacon et al.,
2014; Marocco et al., 2017). Wide stellar companions have been used to
calibrate the ages of main sequence stars (Garcés et al., 2011; Fouesneau et
al., 2019), as well as spectroscopic metallicity relations of M dwarfs (Rojas-
Ayala et al., 2010; Mann et al., 2013, 2014). The use of benchmarks even
extends to groups of related objects, like young moving groups, which are co-
evolving and share common chemical and kinematic histories (e.g., Zuckerman &
Song, 2004).

In addition to serving as benchmarks for the calibration of atmospheric
models, brown dwarf companions present the opportunity to investigate the
formation pathways of substellar objects. Stellar abundance ratios, such as
C/O, have been proposed as a tracer of formation mechanisms (e.g., Mordasini
et al., 2016; Espinoza et al., 2017). If the C/O ratios between the host and
its companion are within statistical agreement it is likely they formed
simultaneously through direct collapse, whereas disagreement in C/O ratio may
be indicative of a core accretion scenario for the substellar companion (Öberg
et al., 2011; Madhusudhan, 2012). However, there is evidence that this
assumption may not hold true, as some brown dwarfs which likely formed via
direct collapse appear to have elevated C/O ratios (Calamari et al., 2022).
Recent works have begun using atmospheric retrievals of brown dwarfs to derive
abundances, temperature-pressure profiles, and other fundamental parameters
(e.g., Gonzales et al., 2020; Burningham et al., 2021; Zalesky et al., 2022;
Lueber et al., 2022; Vos et al., 2023). For benchmark systems, the ability to
compare the derived abundances with those of the host star can ground the
analysis of brown dwarf atmospheric chemistry with external empirical
information (Calamari et al., 2022; Wang et al., 2022a; Hoch et al., 2023;
Gaarn et al., 2023).

In this paper, we present the discovery of 89 low mass star and brown dwarf
(together referred to as ultracool dwarfs) benchmark systems, with spectral
types between M7 and T9, identified through the citizen science project,
Backyard Worlds: Planet 9. The paper is organized as follows: In Section 2 we
describe the citizen science project and the methods of discovery. In Section
3 we discuss the results of our chance alignment probabilities. In Section 4
we describe the acquisition of spectra and archival data. The results of our
analysis are presented in Section 5. Section 6 discusses the different type of
benchmark systems and their distinct uses. We discuss the properties of our
sample in Section 7 and place them into context with the literature. Our
conclusions are presented in Section 8. We include a discussion of noteworthy
and interesting systems in Appendix A.

##  2 Backyard Worlds: Planet 9 Citizen Science Candidate Selection

###  2.1 Selection method 1: Zooniverse Search

Backyard Worlds: Planet 9 (hereafter referred to as BYW) is a citizen science
project hosted on the Zooniverse platform with the goal of identifying brown
dwarfs missed in previous searches, as well as objects within the Solar System
(such as the hypothesised Planet 9) (Kuchner et al., 2017). In the BYW
project, citizen scientists visually inspect an animated ’flipbook’ of time-
resolved 704″ x 704″ coadded unWISE images (for more information see Meisner
et al. 2017a, b, 2018). Because of their red WISE W1–W2 (3.4 $\mu$m – 4.6
$\mu$m) color and relatively fast proper motions, brown dwarfs and other high
proper motion stars stand out in BYW’s animated flipbooks. This platform has
led to many new discoveries by citizen scientists (e.g., Faherty et al., 2021;
Meisner et al., 2020a; Schneider et al., 2020; Jalowiczor et al., 2021;
Gramaize et al., 2022).

Through their searches, citizen scientists can visually observe that their
candidate brown dwarf shares a similar proper motion to another nearby object
within the field of view, and flag their submission as a possible co-mover. In
addition to just using the default flipbooks presented by BYW, many volunteers
also make use of the citizen science developed visualization tool ”Wiseview”
(Caselden et al., 2018), which allows for custom settings (such as FOV,
contrast, speed of ”flipbook”, etc.). In order to create our candidate list,
we first collected all user submissions which were flagged as “comovers”. Each
candidate pair was then visually inspected for authenticity. Using the
CatWISE2020 catalog (Marocco et al., 2021), we gathered the W1 and W2
magnitudes, as well as the proper motions for each candidate brown dwarf
(hereafter referred to as the secondary component of the candidate binary
system). Proper motions, parallaxes, and photometry of all primary components
were taken from the Gaia eDR3 catalog (Gaia Collaboration et al., 2021a),
which was the latest data release at the time of the analysis. We subsequently
used Gaia DR3 values when they became available. Adopting the parallax of the
primary component for the secondary, we plotted each secondary on a W1–W2 vs
MW2 color-magnitude diagram (CMD) as a method to remove any obvious chance
alignments. This was a visual check to eliminate objects that had a
$\sim$3$\sigma$ difference between the spectrophotmetric distance of the
secondary and the parallax of the primary. For example, an object with the
W1–W2 color of a late-type T dwarf but with the absolute W2 magnitude of an M
dwarf would be flagged as an unlikely companion candidate. Seventy-four
candidate systems resulted from this search.

###  2.2 Selection method 2: Gaia to CatWISE Direct Search

The BYW project has been enormously successful at finding a large number of
new brown dwarfs, including co-moving systems. However, all of the co-moving
systems discovered thus-far through BYW have been serendipitous. To supplement
the discoveries already made through BYW, we conducted a dedicated search to
identify candidate proper motion companions using CatWISE2020 (Marocco et al.,
2021) and Gaia DR2 (Gaia Collaboration et al., 2018), which was the latest
Gaia data release at the time of this analysis. All values were updated with
those from Gaia DR3 as those became avaialable.

This search was done via a catalog cross match between CatWISE2020 for the
brown dwarf secondary components, and Gaia DR2 limited to 100 pc (which
included 700,055 sources). CatWISE2020 provides proper motion values for all
sources with average astrometric uncertainties on $\mu_{ra}$ and $\mu_{dec}$
of 10 mas yr-1 (30 mas yr-1) for sources with W1 magnitudes $\sim$12.5 mag
($\sim$15.5 mag; Marocco et al. 2021). Using a Python script, we queried the
CatWISE2020 catalog in a 20 arcminute radius around any Gaia DR2 source with a
parallax $>$ 10 mas and required a $<$3$\sigma$ proper motion component match
or:

$\left|\mu_{\alpha_{Gaia}}-\mu_{\alpha_{Catwise}}\right|<3*\sqrt{\sigma_{\mu_{\alpha_{Gaia}}}^{2}+\sigma_{\mu_{\alpha_{Catwise}}}^{2}}$
(1)

$\left|\mu_{\delta_{Gaia}}-\mu_{\delta_{Catwise}}\right|<3*\sqrt{\sigma_{\mu_{\delta_{Gaia}}}^{2}+\sigma_{\mu_{\delta_{Catwise}}}^{2}}$
(2)

For sources that passed this criterion, we used the ‘ab$\\_$flags’ and
‘cc$\\_$flags’ columns in the CatWISE2020 catalog to filter out as many
artifacts and false detections as possible. These flags contain a series of
letters, corresponding to each of the WISE bands. Each letter represents a
possible artifact: ‘D’ for diffraction spike, ‘H’ for scattered-light halo,
‘O’ for optical ghost, or ‘P’ for charge persistence. If the source is
believed to be real, but contaminated by an artifact, the same letters are
used but in lowercase form. We removed any CatWISE source which contained a
capital letter flagging either the W1 or W2 bands. As CatWISE2020 does not
contain W3 or W4 photometry, we made no cuts to remove artifacts in those
bands. We also did not remove any sources with lowercase letters to avoid
cutting any potentially real sources.

What remained afterwards was a list of $\sim$40,000 candidate co-moving pairs
to a known Gaia DR2 star. We then applied a color cut so that only sources
with a W1–W2 color $>$ 0.2 mag remained, restricting the candidate secondaries
to spectral types later than $\sim$M7 (Skrzypek et al., 2016a). Next, we
estimated the projected physical separation of the system by combining the
distance of the primary with the separation between Gaia and CatWISE sources.
Based on the separation distribution of known wide co-moving systems with a
low mass companion (e.g Faherty et al. 2010a), we eliminated any system with a
projected separation $>$ 30,000 AU in order to minimize chance alignments.
Since pairs where both components have a detection in Gaia have been the focus
of other searches for co-moving pairs (cf. El-Badry et al. 2021; Zhou et al.
2023), we were only interested in pairs with a CatWISE-only detected
companion. We therefore eliminated all pairs in which both objects had entries
in Gaia. However, it should be noted that no Gaia-Gaia or $>$30,000 AU pairs
were removed from the sample selected in Section 2.1.

Our remaining list consisted of 5,963 candidate pairs. To verify that each
system was authentic (both objects being real, and actual motion seems to
match) we only retained pairs which had a total proper motion $\geq$100
mas/yr, resulting in 3,135 candidates. In order to parse through the large
number of candidates, we used the machine learning algorithm described in
Marocco et al. (2019) to help assign a likelihood to each candidate ultracool
companion. This machine learning algorithm, based on the python code XGCBoost
(Chen & Guestrin, 2016), was trained on red and faint brown dwarfs from the
literature, and assigns a “score” on how likely a source in CatWISE2020 is to
be real (for further details, see Marocco et al. 2019). Retaining only sources
which had a score $>$90$\%$, the list of candidate pairs was reduced to 1,278
pairs. Sixty percent of this final list has been visually inspected, and 58
candidates have been identified, with 19 of those overlapping with the
candidates identified in Section 2.1. The remaining sources inspected
contained either spurious detections, known co-moving pairs, or companions
which did not have visually matching proper motions.

##  3 Chance Alignment Probabilities

In order to assess the probability of each system being a real co-moving pair
and not a chance alignment, we used a modified version of BANYAN $\Sigma$,
called CoMover (Gagné et al., 2021). This code uses the proper motion
components, parallax or distance estimate, heliocentric radial velocity (where
available), and sky position – all with uncertainties – for both the host star
and its potential companion. Utilizing a six dimensional multivariate Gaussian
in Galactic coordinates and space velocities, a single spatial-kinematic model
is constructed from the input of the host star. The observables of the
companion are then compared to the host’s model, as well as to the field-star
model of Gagné et al. (2018). The code also uses Bayes’ Theorem to marginalize
over any of the companion’s missing data with the analytical integral
solutions from Gagné et al. (2018). The output of CoMover is a probability
that the host and companion are a real, co-moving system.

###  3.1 High Probability Systems

Due to the use of Bayes’ Theorem in CoMover, the output for a low probability
system ($\lessapprox$70$\%$) is inherently unstable. Multiple runs of the code
using the same observables can result in different probabilities. However,
this has not been observed to produce a false positive, rather the different
resultant percentages are all still low. We observed that outputs which are
above 90$\%$ appear to converge on the same probability. We therefore make a
conservative cutoff at 90% as what constitutes a “high probability” co-moving
system.

Seventy-nine percent of the candidate pairs identified in this paper received
a co-moving probability above this threshold, resulting in 89 high probability
systems. A summary of these high likelihood systems is listed in Table 1, with
their respective photometry shown in Table 2. Hosts and companions are shown
on CMDs (using the host star’s parallax for the companions) in Figure 1.

###  3.2 Low Probability Systems

Twenty-four of the candidate pairs considered received a co-moving probability
below our 90$\%$ threshold using CoMover. These twenty-four systems, which
were originally identified as potential co-moving pairs, thus have a high
likelihood of being chance alignments, and we therefore do not consider them
in the discussion which follows. Three of the ultracool dwarfs in these low
likelihood pairs had been followed up spectroscopically, and are presented in
Figure 2. A summary of the low likelihood pairs is shown in Table 3.

##  4 Data Acquisition

###  4.1 Spectroscopic Observations

We have obtained follow-up spectroscopy of 34 of the potential companions
identified in Section 2. Thirty-three candidate hosts were also observed. A
summary of the observation dates and instruments used is shown in Table 4.

####  4.1.1 IRTF/SpeX

We observed twenty-two objects with NASA’s Infrared Telescope Facility (IRTF)
located on MaunaKea, Hawaii. Using the SpeX instrument (Rayner et al., 2003)
we obtained near infrared spectra over the 0.8–2.5 $\mu$m wavelength range.
All objects were observed with SpeX in prism mode with a 0$\farcs$8 wide slit,
providing a resolving power of $\approx$100-500. Additionally, one target was
observed using SpeX in the short-wavelength cross-dispersed mode (SXD), with a
wavelength coverage of 0.8–2.42 $\mu$m with a resolving power of R$\sim$2000\.
After each target, a standard A0 star was observed for telluric correction and
flux calibration. The data was reduced with Spextool (Cushing et al., 2004)
using the telluric and flux calibration techniques of Vacca et al. (2003).

####  4.1.2 Keck/NIRES

Ten objects were observed using the Near Infrared Echellette Spectrometer
(NIRES; Wilson et al. 2004) at the W. M. Keck II telescope, located on
Maunakea, Hawaii. Near infrared spectra (0.94–2.45 $\mu$m) were obtained at a
resolving power of $\approx$2700\. Data were acquired during five nights on 27
October 2018, 14 February 2019, 24 February 2021, 21 July 2021, and 11 June
2022 (UT). Conditions were generally clear with 0$\farcs$5 to 1$\farcs$0
seeing, with the exception of clouds on 27 October 2018 and 21 July 2021. Each
source was acquired with the K-band imaging channel of NIRES and observed with
the 0$\farcs$55 $\times$ 18$\arcsec$ slit aligned to parallactic, and 4-8
exposures were acquired in an ABBA 10$\arcsec$ dither sequence with
integration times of 10 s to 300 s depending on source brightness and
conditions. We observed an A0 V star before or after each science target for
telluric absorption and flux calibration, and dome-reflected flat-field and
arc lamps for order tracing, pixel calibration, and initial wavelength
calibration. All data were reduced using a modified version of SpeXtool
(Cushing et al., 2004; Vacca et al., 2003) following standard reduction
procedures.

####  4.1.3 Lick/Kast

Twenty objects were observed using the Kast double spectrograph (Miller &
Stone, 1994) on the Shane 3m Telescope at the Lick Observatory located on
Mount Hamilton, California. Data were acquired over six nights on 17 May 2021,
12 December 2021, 29 September 2022, 31 October 2022, 1 June 2023, and 12
August 2023 (UT) in conditions that ranged from clear to thin cirrus and
seeing of 1$\farcs$1–1$\farcs$6\. All sources were observed using the 600/7500
grism and 1$\farcs$0- or 1$\farcs$5-wide slit, providing 6000–9000 Å
wavelength coverage at an average resolution of $\lambda/\Delta\lambda$ =
1900–2800. Individual exposures ranged from 30 s to 1200 s depending on source
brightness, and were obtained in pairs to reduce cosmic ray contamination. For
each source, we observed a nearby G2 V star for telluric absorption
correction, and on each night we observed a flux calibrator drawn from Hamuy
et al. (1992, 1994). Data were reduced using the kastredux
package111kastredux: https://github.com/aburgasser/kastredux. using the
baseline parameters for batch reduction.

####  4.1.4 SOAR/TripleSpec4

Three objects were observed with the SOAR Telescope at the Cerro Tololo Inter-
American Observatory (CTIO), using the Astronomy Research using the Cornell
Infra Red Imaging Spectrograph (ArcoIRIS; Schlawin et al. 2014) otherwise
known as TripleSpec4. Near infrared spectra (0.8–2.4 $\mu$m) were obtained
across 6 simultaneously observed cross-dispersed orders covering the 0.8 - 2.4
$\mu$m range, with a resolving power of $\approx$3500\. Science exposures were
taken at two different nod positions along the slit, which has a fixed width
of 1.1$\arcsec$. After each science target, A0V stars were observed in order
to execute telluric corrections. Data were obtained on 2021 October 23 under
stable and clear conditions. Reductions were performed using a modified
version of the SpeXtool reduction package (Cushing et al. 2004; Vacca et al.
2003).

####  4.1.5 SALT/RSS

Seven objects were observed with the Robert Stobie Spectrograph (RSS)(Burgh et
al., 2003; Kobulnicky et al., 2003), located on the Southern African Large
Telescope (SALT)(Buckley et al., 2006). The long slit mode of the spectrograph
was used, with the PG0900 grating at a 20o angle. This gives coverage over 3
wavelength ranges: 6033-7028 Å, 7079-8045 Å, and 8091-9023 Å. Neon arc lines
were observed immediately after each target, which was used to wavelength
calibrate the observatory-provided pre-processed data (included gain and cross
talk correction, as well as overscan subtraction). The target’s spectrum was
then flux calibrated using the Hamuy et al. (1994) standard EG21, which was
acquired on 2023 January 24 UT using the same setup.

###  4.2 Catalog Photometry, Astrometry, and Abundances

We performed a literature search for any existing near-infrared photometry in
the Two Micron All Sky Survey (2MASS) (Skrutskie et al., 2006), various UKIRT
surveys (e.g., UKIDSS; Lawrence et al. 2007, UHS; Dye et al. 2018), and from
the VISTA science archive (Cross et al., 2012a) for both host and companion
objects.

For all candidate host stars we gathered proper motions, parallaxes, radial
velocities, and rotational broadening ($v$sin$i$) using Gaia DR3 (Gaia
Collaboration et al., 2023a). Additionally, Gaia DR3 contains abundances and
metallicity for select sources, which were also gathered if present. Further
literature searches were performed for any previously reported abundances or
metallicity. For sources which had metallicity measurements available in the
literature, no corrections were added to account for systematic offsets
between different studies.

###  4.3 TESS Light Curves

As the characterization of brown dwarfs is heavily age dependent, being able
to derive an estimate for the host star age in a benchmark system is crucial.
Gyrochronology, the use of a star’s rotation rate to constrain its age (e.g.,
Barnes, 2003, 2007), has been shown to provide reliable age constraints for
solar and later type stars for ages of a few hundred Myr out to field age.
With new large area variability surveys, such as NASA’s Transiting Exoplanet
Survey Satellite (TESS; Ricker et al. 2015), obtaining rotation rates for
large numbers of stars has become much easier, increasing the power and
usefulness of gyrochronology as an age diagnostic tool.

Host sources were checked to see if any had been observed by TESS with either
30 minute or 2 minute cadences. If observed, a light curve was generated using
simple aperture photometry and background subtraction. Each light curve was
then run through a LombScargle periodogram (Lomb, 1976; Scargle, 1982),
resulting in a best fit rotation period (see Popinchalk et al. 2023a for
further details regarding the process used).

##  5 Analysis

###  5.1 Spectral types

####  5.1.1 Primary Components

For candidate primaries which were observed spectroscopically, spectral types
were assigned by overplotting either infrared spectral standards (using the
standards of Kirkpatrick et al. (2010a) for types M7 and later, and sources
from the IRTF/SpeX spectral library for sources earlier than M7) or optical
standards (using standards from Kirkpatrick et al. (1991) for spectral types
K5 through M9, and Kesseli et al. (2017) for earlier types) on top of the
observed data, normalizing all spectra and standards to the same wavelength.
In the optical, spectral types were assigned via a visual match to the entire
wavelength coverage, whereas spectral types in the near-infrared were assigned
based on a visual match of the J band portion of the spectrum following the
technique of Kirkpatrick et al. (2010a). The primary near-infrared and optical
spectra are shown in Figures 3 and 4 respectively.

For candidate primaries which did not have spectral types or phototypes in the
literature, and were not observed spectroscopically, we used the Gaia parallax
and photometry along with the absolute magnitude relations of Kiman et al.
(2019) to estimate spectral types. If the candidate primary lacked a Gaia
parallax, its spectral type was estimated from the spectral type versus color
relations from Kirkpatrick et al. (2021).

####  5.1.2 Secondary Components

For the secondaries, spectral types were assigned by overplotting the infrared
spectral standards from Burgasser et al. (2006) and Kirkpatrick et al. (2010a)
on top of the observed spectra and visually comparing the peak of the J band.
All secondary spectra are shown in Figures 5 (near-infrared spectra) and 6
(optical spectra).

For candidate secondaries with no available spectra, we estimated the spectral
type by adopting the host star’s parallax and using the spectral type versus
$M_{W2}$ relation from Kirkpatrick et al. (2021). If the candidate secondary
was a companion to an object which lacked a parallax, the secondary’s spectral
type was estimated with the spectral type versus color relations from
Kirkpatrick et al. (2021) For candidate secondaries which had their own Gaia
astrometry and lacked spectra, spectral types were estimated using the
photometric relations from Kiman et al. (2019).

###  5.2 Low Probability Systems

While the low probability systems presented in Section 3.2 have a high
likelihood of being chance alignments, due to large projected physical
separations and differences in proper motions, 6 systems still receive
comoving probabilities $>$60$\%$. Of these, 3 systems have probabilities
85$\%$: G 1-44 + 2MASS J0102+0355 (88.6$\%$), WDJ1324-1902 + CW1325-1907
(87.5$\%$), and 2MASS J2308-5052 + CW2308-5052 (87.2$\%$). These systems,
while not above the “high probability” threshold, still have elevated
probabilities of being associated with one another. These candidate companions
have proper motion errors which range from $\sim$34-70″, making it difficult
to confirm their comoving nature. A re-analysis with higher accuracy proper
motions is required in order to further investigate the candidacy of these
systems.

Three of the ultracool dwarfs in the low probability sample were followed up
spectroscopically:

CW1111-4454 Comparing the near-IR spectrum of CW1111-4454 to that of the
spectral standard in Figure 2, we place CW1111-4454 as spectral type L2 (sl.
blue), due to the slightly supressed K-band of CW1111-4454 which gives it a
slightly bluer spectral morphology.

CW0230+2405 We place CW0230+2405 as spectral type L4 (blue) through comparison
with the L4 spectral standard as shown in Figure 2. Both the H- and K- bands
of CW0230-2405 appear to be suppressed compared to the standard.

CW0358+5624 CW0358+5624 is best fit as spectral type L7, as shown in Figure 2.

###  5.3 Host RUWE

The Gaia re-normalized Unit Weight Error (RUWE) (Lindegren et al., 2018) is
related to how ‘well behaved’ the astrometric and photometric solution is for
a source when compared to the proper motion and parallactic motion expected
for a single source, while taking into account the source’s magnitude and
color. A ‘perfectly well behaved’ RUWE value would be expected to be
$\sim$1.0. Deviations from this value are caused by non-linear motions of the
photocenter of the source in the Gaia images, as well as unusual colors (e.g.
as a product of variability). Therefore, an elevated RUWE value can be
attributed to any potential variability of the source, as well as motion
caused by a binary orbit.

Several studies have been performed on testing the RUWE with known binary
systems, and it has been shown that the RUWE strongly correlates with
unresolved binary companions (Stassun & Torres, 2021; Belokurov et al., 2020).
Nominally, RUWE values $\gtrsim$1.4 are considered to be poorly behaved,
although Stassun & Torres (2021) have shown that even RUWE values just
slightly larger than 1.0 can be indicative of an unresolved companion.

We gathered the RUWE values for all host stars, and find that 20 sources have
a RUWE 1.4, as shown in Table 11.

###  5.4 Gravity Indices

Young ultracool dwarfs, which have not yet fully contracted, should have lower
surface gravities than their field counterparts. Thus, being able to determine
the surface gravity of an ultracool dwarf can provide insight into the
object’s age.

For objects between M6 - L7 which had near infrared spectra, we calculated the
FeH-z, VO-z, KI-J, and H-cont indices of Allers & Liu (2013), which have been
found to be sensitive to surface gravity. These indices are not appropriate
for spectral types earlier than M6 or later than L7, so these spectra were
left out of the analysis. Following the approach of Allers & Liu (2013), we
score each of the indices based on spectral type, and assign an overall
classification for the object of field gravity (FLD-G), intermediate gravity
(INT-G), and very low gravity (VL-G). We find 7 companions which receive an
INT-G classification, and 19 which are classified as FLD-G. These results are
shown in Table 5.

###  5.5 Spectral Binary Indices

For ultracool dwarfs in tight binary pairs which are unresolved, the combined
light from the two objects can result in an observed spectrum resembling that
of a single, typical ultracool dwarf (Burgasser et al., 2010; Bardalez
Gagliuffi et al., 2014). These systems are referred to as spectral binaries.
If the two components of the spectral binary are different enough from one
another (such as an L- and T-dwarf), the observed spectrum is likely to
contain peculiarities which arise from the different spectral features between
the two objects.

In order to investigate if any of the observed spectra of the candidate
companions in our sample are indicative of spectral binarity, we followed the
methods of Bardalez Gagliuffi et al. (2014) for spectral types M7–L7 and
Burgasser et al. (2010) for spectral types L8–T5. Table 6 lists the indices
calculated for each object, including the strength of spectral binary
candidacy based on the definitions of Burgasser et al. (2010) or Bardalez
Gagliuffi et al. (2014) (weak, strong, or rejected). Of the 29 companions
included in this analysis, 2 objects were found to be strong spectral binary
candidates, 7 objects were found to be weak spectral binary candidates, and 20
objects were rejected as spectral binaries. Futher discussion of these SB
candidates can be found in Section 7.6.2.

###  5.6 Rotation Periods

Seventy-three of the host stars in our sample were bright enough to be
detected with TESS. Of these seventy-three, we were able to confidently
measure a rotation period for $\sim$14$\%$ (10 sources). These sources are
listed in Table 7. A sample light curve and periodigram from our workflow is
shown in Figure 7. We placed these stars on a color-period diagram along with
stars from the Tucana-Horologium (40 Myr, Popinchalk et al. 2023b), Pleiades
(120 Myr, Rebull et al. 2016) and Praesepe (650 Myr, Douglas et al. 2017)
clusters as age benchmarks, shown in Figure 8.

###  5.7 BANYAN $\Sigma$

Each candidate system was checked for young moving group memberships
probabilities using BANYAN $\Sigma$ (Gagné et al., 2018). BANYAN $\Sigma$ is a
Bayesian algorithm which uses the spatial and velocity coordinates of a source
in order to identify potential membership to young moving groups (YMG). For
each source, BANYAN $\Sigma$ uses its sky coordinates, proper motions, radial
velocity, parallax, and uncertainties, comparing them in full 6D configuration
space (XYZUVW) to 27 known young associations (see Gagné et al. 2018 for a
detailed explanation). BANYAN $\Sigma$ also allows for the possibility of
using only sky coordinates and proper motions if no other measurements exist.

The output of BANYAN $\Sigma$ is a list of probabilities for membership to
each of the 27 YMGs, along with a probability of belonging to the ”field”
population. The ages of the YMGs considered in BANYAN $\Sigma$ range from
$\sim$1-800 Myr. As BANYAN $\Sigma$ only contains 27 YMGs, a source receiving
a high probability of belonging to the field can only be said to not likely
belong to any of those 27 groups. There still exists the possibility the
source may still belong to one of the other known YMGs which are not included
in the analysis of BANYAN $\Sigma$.

For each candidate system in our sample, we ran all available kinematics for
both host and companion through BANYAN $\Sigma$ in order to check for any
potential YMG memberships. Systems which received a likely membership are
discussed further in Appendix A.

##  6 Discussion of Benchmark Systems

Benchmark systems are extremely useful, as they provide a means of gaining
insight on fundamental properties like distances, chemical abundances, and age
from a co-moving system where one object is very well characterized and can
inform on the other. The sample of ultracool benchmark systems presented in
this paper comprise a wide variety of host spectral types, each of them being
able to benchmark different aspects of their companion’s properties. Below we
describe the different “populations” of brown dwarf benchmarks in our sample,
as well as the vital role each can play.

###  6.1 Age Benchmarks

Twenty-four of the host stars presented in this sample provide a means of
reliably constraining an age range for their ultracool dwarf companions. These
ages are listed in Table 8. This added information makes it possible to break
the mass-temperature-age degeneracy of substellar mass objects, which will be
utilized in future works aiming to derive fundamental properties of the BYW
benchmark sample.

Six of the hosts (or at least one component of the host system) in our sample
are white dwarfs: WDJ060159.98-462534.40, LP 369-22, UCAC4 328-061594,
WDJ002027.83-153546.50, WDJ011001.84-592642.22, and WDJ212231.86+660043.69.
Using the python code wdwarfdate (Kiman et al., 2022), we input the effective
temperature and log(g) of the white dwarf from Gentile Fusillo et al. (2021),
and get the Bayesian estimated white dwarf mass, progenitor star mass, white
dwarf cooling age, and the total system age (including the estimated lifetime
of the progenitor star). The ages obtained for these sources range from
$\sim$0.42-10.91 Gyr. Kiman et al. (2022) state that the cooling tracks
employed by wdwarfdate assume a C/O core, and therefore the ideal mass range
for this method is between 0.45-1.1 $M_{\sun}$. Likewise, Kiman et al. (2022)
point out that white dwarfs with masses $<$0.45 $M_{\sun}$ and $>$1.1
$M_{\sun}$ may have potentially formed as a result of binary
mergers/evolution. Results for white dwarfs outside of this ideal range should
be treated with extreme caution. Two of our white dwarf hosts lie outside of
this range: WDJ011001.84-592642.22 (M = 0.43 $M_{\sun}$) and UCAC4 328-061594
(M=1.19 $M_{\sun}$); these sources require further study to more accurately
constrain their ages. For more information on the limitations and errors of
wdwarfdate, refer to Kiman et al. (2022).

While we do believe a seventh host star to be a white dwarf, its photometry is
contaminated, preventing us from properly estimating its age. Further
observations of this source, however, should provide enough information to be
able to constrain its age.

Nine of our hosts are main sequence stars with age estimates in the
literature, spanning $\sim$0.1-4.2 Gyr. Seven of these age estimates utilize
the equivalent width of the CaII infrared triplet excess from Žerjal et al.
(2017), noting that this approach only allows us to place loose age
constraints on these systems (i.e. 0.1–1 Gyr or $>$1 Gyr).

For ten of our host stars, we were able to obtain a light curve from TESS
observations. An additional 3 of our host stars have reported rotational
periods in the literature which are not derived from TESS. We compare these
sources to confirmed members of the Pleiades (120 Myr) and Praesepe (650 Myr)
groups on a color-period diagram (see Figure 8). From comparing the rotation
periods of our objects with that of these benchmark clusters, we estimate that
all but two sources appear to be $<$1 Gyr, with 5 appearing to be potentially
younger than the $\sim$650 Myr Praesepe cluster, based on their rotation
periods. These ages should be considered estimates based on comparison to
benchmark populations and not derived from robust gyrochrones. Two of these
hosts stars have a companion which received an INT-G gravity score, while the
companions for the rest either lack spectra or lack clear gravity sensitive
indices for their spectral type.

One of our host stars, GJ 900, is a potential member of the nearby moving
group Carina-Near, which has an age of 200$\pm$50 Myr (Zuckerman et al.,
2006).

###  6.2 Compositional Benchmarks

One of the most difficult properties to measure for an ultracool dwarf is
chemical composition (e.g. metallicity, C/O ratio, Mg/Si ratio). Chemistry is
challenging to derive for most stars given that high resolution spectroscopy
is required and solar type stars are the optimal spectral type for
measurements. The intrinsic faintness of ultracool dwarfs makes obtaining the
high-resolution spectra necessary for these measurements difficult for many
objects. To date, only $\sim$50 wide, resolved ultracool companions with a
spectral type of L or later to solar-type analogs (FGK stars) are known
(Faherty et al., 2010b; Deacon et al., 2014, Calamari et al. in prep). In this
paper, we present an additional 21 ultracool companions to FGK stars,
consisting of 3 late M dwarfs, 14 L dwarfs, and 4 T dwarfs, expanding the
known widely separated companions to FGK stars by $\sim$42%. Each of the host
stars has the potential to be chemically characterized and then the secondary
can have its observational features calibrated. Ten of the host systems have
available abundances in the literature thanks to large area surveys such as
LAMOST or Gaia DR3. These systems, and their published abundances are listed
in Table 9.

There are two ways the ultracool dwarf compositional benchmarks can be
utilized: (1) with a large enough sample of L and T dwarfs co-moving with an
FGK star, we can correlate the host star compositions with observational
properties of the secondary. For instance, we can look for metallicity
sensitive features across the near infrared and map that by the metallicity
measurements of the primary. (2) Utilizing the well defined properties of the
host star, we can constrain the retrieved properties of the companion (such as
metallicity and abundances), testing ultracool dwarf and exoplanet atmospheric
models.

Fifty-five of our host stars have metallicities reported in the literature,
shown in Table 9. The range of metallicites (both [Fe/H] and [M/H]) for these
stars are from -2.03 to 0.55 dex, with fifteen having a metallicity $<$ -0.5
dex. Adopting the host’s metallicity for their ultracool companions, the wide
range of metallicities in this sample can help address the role it plays in
shaping the spectral energy distributions of cold atmospheres.

Two of the hosts in our BYW benchmark sample allow us to derive C/O abundance
ratios from their reported [C/H] and [O/H] metallicites. The range of these
abundances are from C/O=0.44 to 0.79. While potentially useful in studying
ultracool dwarf formation pathways, comparing the C/O ratios between a host
and its ultracool companion can serve another important purpose. Calamari et
al. (2022) point out a trend in the retrieved C/O values in late-T dwarfs to
be supersolar. While this could be hinting at formation, Calamari et al.
(2022) postulate that this trend is likely due to unconstrained CO and CO2
abundances, as well as potentially misunderstood or unconstrained condensate
chemistry. Many of these late-T dwarfs with retrieved C/O values do not have a
companion with which to compare their abundances, however. While the retrieval
work is deferred to future papers, the sample presented in this paper allows
for more direct comparisons between the atmospheric abundances retrieved in
ultracool dwarfs with their host, hopefully beginning to shed more light on
the observed trends.

Three of the hosts in our BYW benchmark sample allow us to derive Mg/Si
abundance ratios from their reported [Mg/H] and [Si/H] metallicities. The
Mg/Si abundances for these sources ranges between Mg/Si=0.68 to 1.4. In many
ultracool atmospheres, the temperatures become cool enough for the formation
of silicate and iron clouds (Lodders & Fegley, 2006). From models of brown
dwarf atmospheres, these clouds become especially prominent in the mid-late L
dwarfs and across the L-T transition, until they sink below the photosphere in
T-dwarfs (Lodders & Fegley, 2006). Observations agree nicely with this general
picture, with dust clouds being found to peak between the L4 and L6 spectral
types, sedimenting at spectral type $\sim$ L8 (Suárez & Metchev, 2022, 2023)
The species of silicate clouds which can form in these atmospheres has direct
implications in modelling the observed atmosphere, as different silicate
clouds will sequester different amounts of oxygen as well as shape the
temperature - pressure profile. Atmospheres with a high Mg/Si (magnesium rich)
are more likely to form forsterite (Mg2SiO4) and enstatite (MgSiO3) clouds,
however for low Mg/Si atmospheres (silicon rich), enstatite and quartz (SiO2)
are the expected species (Calamari et al. in prep). Assuming both host and
companion formed with the same elemental abundances, knowing the Mg/Si ratio
of the host allows for a prediction of the clouds expected in the companion,
directly testing the retrieved cloud species. With a large enough sample of
ultracool dwarf companions which can have their atmospheric chemistry anchored
via their host’s composition, it may be possible to begin constraining
theories of cloud formation and sedimentation in both ultracool dwarfs and
exoplanets. The host stars in this sample provide a large range of Mg/Si
abundances, which can be utilized to directly test varying cloud formation
models in their ultracool companions.

##  7 BYW Benchmark Sample in Context with Literature Benchmark Sample

Within this section we summarize several critical properties of the new co-
movers in the context of the literature sample of benchmarks.

###  7.1 Distance distribution of The BYW Benchmark sample

The overall distances covered by the full BYW benchmark sample ranges from
$\sim$20 pc out to $\sim$170 pc, with the peak of the distribution around 40
to 60 pc. For the ultracool companions which are classified as late M’s, the
distance distribution is from 45 pc to 174 pc. While the majority of these
companions, which were discovered using the method described in Section 2.1,
are also detected by Gaia along with their hosts, four companions have no Gaia
counterpart: CW1240+4605, CW1325+0223, CW0159+1055, CW1031+1237. The distances
for these objects (adopting the Gaia parallax of their host) are
124.0${}_{-4.8}^{+5.2}$ pc, 119.0${}_{-0.4}^{+0.4}$ pc,
111.0${}_{-0.6}^{+0.6}$ pc, and 99.0${}_{-0.2}^{+0.2}$ pc respectively. The
Gaia Catalog of Nearby Stars (GCNS, Gaia Collaboration et al. 2021b) is 95$\%$
complete out to 100 pc down to spectral type M8, so it is not unexpected for
objects with spectral type M9 beyond 100 pc to be missing astrometry.
Additionally, two companions, both spectral type M9, have a Gaia DR3 catalog
entry with no astrometry indicating that they were likely at the edge of
detection for the mission.

The distance of the L-dwarf type companions in the BYW benchmark sample ranges
from 30 pc to 146 pc, with the distribution peaking around $\sim$55-60 pc. Two
of the companions which had been identified using the method described in
Section 2.1 have entries in Gaia DR3: CW2003$-$1422 (spectral type L0 at 56.99
pc) and CW2029$-$7910 (spectral type L1 (blue) at 50.95 pc). One companion,
CW1303+5127 (spectral type L2), has a Gaia DR3 catalog entry with no
astrometry.

The distances of the T-dwarf type companions in the BYW benchmark sample
ranges from 17 pc to 57 pc, with the distribution peaking around $\sim$50 pc.

###  7.2 Separation of BYW Benchmark Systems

Characterizing the demographics of ultracool companions is important for
testing and constraining different formation mechanisms (Metchev &
Hillenbrand, 2009; Bowler et al., 2014; Brandt et al., 2014). While the
occurrence rates and characteristics of widely separated ($>$ 100 AU)
ultracool companions have been well studied, only $\sim$46 ultracool
companions have been discovered with separations $\geq$1000 AU (see Chinchilla
et al. 2020). In this paper, we add an additional seventy-two systems to the
list of systems with separations $\geq$1000 AU, almost tripling the current
census.

Presently, some of the widest known confirmed ultracool companions have been
observed to have projected separations $\sim$15,000 AU (Zhang et al., 2010;
Day-Jones et al., 2011). The BYW benchmark sample includes twelve systems with
companions at separations $>$15,000 AU. Interestingly, seven out of the twelve
of these widest systems have T dwarf companions. The formation of these
systems is difficult to explain with current models, as in situ formation of
binary systems via direct collapse and fragmentation of the molecular cloud
seems to only form pairs out to $\sim$10,000 AU (Chen et al., 2013). One
possible explanation for their formation is that these wide systems are
actually of higher order multiplicity, where the interactions between
components result in an extremely wide orbit (Reipurth & Mikkola, 2012;
Elliott & Bayo, 2016).

Indeed, we find that five of our wide companions have hosts with multiple
components: CW0915+2547 (estimated T5) with K6 and DZ white dwarf hosts;
CW1959-6443 (estimated T6) with F3IV and M5 dwarf hosts; SDSS 1344+0839 (T0,
identified as a low probability L5+T2 binary) with an M5.5 host which is
likely binary (Gaia RUWE=8.3); CW1625+7749 (L4) with candidate binary
(RUWE=1.6) M2 dwarf host, and CW0407+1909 (estimated T1) with M4 and M4 dwarf
hosts. Additionally, two of the companions, CW1959-6443 and CW0407+1909,
appear overluminous on the WISE CMD, indicating the potential for hidden
binarity in these components as well.

Three of the systems with separations $>$15,000 AU have hosts (or a host
component) which are white dwarfs: CW0915+2547, CW1010-2423, and CW2123+6556.
As the progenitor star undergoes its transformation into a white dwarf, the
mass loss taking place may cause the orbital separation of the host and
companion to increase. Burleigh et al. (2002) showed that this increase should
scale as $M_{ms}$/$M_{wd}$, where $M_{ms}$ is the initial mass of the
progenitor star, and $M_{wd}$ is the final mass of the white dwarf. Therefore
if a system formed with a separation of 5,000 AU, and the progenitor star was
five times as massive as its white dwarf stage, the companion would eventually
migrate out to 25,000 AU. While this migration may in fact lead to the pair
dissolving as the binding energy becomes too small, because of the lack of a
sudden and violent interaction, the pair could retain a similar proper motion
and distance for a significant time (before drifting apart and being
thermalised by the disk population).

The widest of the systems in the BYW benchmark sample, CW2106+2507, is an
estimated T1 dwarf which was given a 100% probability of co-moving with a K6.
The projected separation of the pair at the Gaia distance of the host is
$\sim$38,500 AU. Neither component shows obvious signs of hidden binarity,
making it unlikely this pair is still gravitationally bound to one another, as
their binding energy would be too small to survive the gravitational
interactions with surrounding objects. Their receiving a high probability of
co-moving with one another hints at them still being associated with one
another in some way, however, and could potentially be indicative of a
dissolving moving group.

###  7.3 Distribution of Host Mass

Host mass estimates were taken from Stassun et al. (2018) and Stassun et al.
(2019) for main sequence hosts, and Gentile Fusillo et al. (2021) for WD
hosts. When no published mass was present, we estimated its mass using the
relations of Pecaut & Mamajek (2013).

The host stars in the BYW benchmark sample cover a wide range of masses, from
$\sim$0.1 $M_{\sun}$ out to $\sim$1.4 $M_{\sun}$. The peak of the distribution
is around 0.2-0.3 $M_{\sun}$, or spectral types between $\sim$M3-M6. Previous
studies have found no discernible difference in the brown dwarf companion
fraction based on host spectral type (Bowler et al. 2015 and references
within), so the observed abundance of mid-M dwarf hosts in our sample is
likely not due to a preference for these spectral types to host ultracool
companions, but is instead due to these spectral types representing the peak
of the spectral type distribution, with mid-M dwarfs comprising $\sim$75$\%$
of the solar neighborhood (Henry et al., 2006).

###  7.4 System Multiplicity

###  7.5 Host Stars

Twenty-one of the new co-moving systems reported here-in have primary stars
with hints of binarity due to large Gaia RUWE values.

In addition to just the single RUWE from Gaia DR3, it is also helpful to
compare the RUWE value from Gaia DR2, as DR3 has a longer baseline of
observations (Lindegren et al., 2021). If a binary system has a long orbital
period, there may not be enough time in the relatively short Gaia observation
window to detect any deviation away from the single star model. However, if
the orbital period was just right, the extra observations between DR2 and DR3
may be enough for the non-single star behavior to show up, causing an increase
in the RUWE. Similarly, the extra observation time could result in the
photocenter wobble averaging out, resulting in a decrease of the RUWE.

We therefore compared the RUWE values from DR2 and DR3 of all the primary
stars of our co-moving systems. Using the nominal RUWE cutoff of $\sim$1.4, we
find that twenty primaries were above this limit in Gaia DR3. These sources,
as well as the hosts which are confirmed binaries, are listed in Table 11.
Looking at the locations of these stars on the Gaia CMD shown in Figure 11, we
find that a large majority of them appear to be overluminous, consistent with
the interpretation of these sources being binaries. A few of the stars with a
high RUWE value appear to be near normal on the Gaia CMD for their spectral
type. While this could be indicating other sources for their anomalous RUWE
values than binarity, this could also be due to a potential hidden companion
faint enough to not contribute much to the total flux of the host, yet bright
enough to effect the objects photocenter.

###  7.6 Companion Multiplicity

####  7.6.1 Overluminous Candidates

In addition to considering the binarity of the host stars, we can also take a
look at the multiplicity of the companions. Ultracool dwarfs in wide binary
pairs have been shown to have higher occurrence rates of binarity (Burgasser
et al., 2005; Faherty et al., 2010a; Law et al., 2010). Looking at our
companions on the WISE CMD in Figure 1, adopting the parallax of their hosts,
we find that seven appear to be overluminous for their spectral type. Six of
the companions that appear overluminous are estimated to have spectral types
T5 or later, and have projected separations from their hosts $>1000$ AU. The
wide separations and low mass ratios lead to very small binding energies for
these systems (see Section 7.7). Assuming the overluminosity of the companions
is due to binarity, the binding energies of these systems could increase and
begin to more closely resemble the binding energies observed in typical binary
systems.

####  7.6.2 Spectral Binary Candidates

Additionally, as shown in Table 6, 9 of the potential companions in the BYW
benchmark sample are candidate spectral binaries.

Strong Candidates Two of our candidate companions are strong spectral binary
candidates: CW0620+3446 (best single fit L4, best SB fit L4+T7) and
CW2029-7910 (best single fit L1 blue, best SB fit sdL1+L2.5). As shown in
Table 11, CW2029-7910’s host star, SCR J2029-7910, is itself a candidate
binary. If the binary nature of both components are confirmed, this would make
the CW2029-7910 + SCR J2029-7910 a widely separated quadruple system.

Weak Candidates Seven of our candidate companions are weak spectral binary
candidates: CW0227+0830 (best single fit L5 blue, best SB fit L5+L6),
CW0627-0028 (best single fit T0 blue, best SB fit L5+T2.5), CW1313+2230 (best
single fit L3, best SB fit L2+T1.5), CW1328-0637 (best single fit L3 blue,
best SB fit L2+L5), CW1439+2600 (best single fit T3, best SB fit T1.5+T4.5),
CW1606-1032 (best single fit L7 blue, best SB fit L4.5+L6.5), and CW2141-0246
(best single fit L4 sl. red, best SB fit L2+L3). Bardalez Gagliuffi et al.
(2014) note that blue L-dwarfs are a large contaminant in the analysis of weak
spectral binary candidates, so blue sources such as CW1328-0637 and
CW1606-1032 are potentially contaminants. The other weak candidates require
further study in order to confirm their binarity.

###  7.7 Binding Energy

The binding energies of our systems were calculated with the equation:

$1.8\times 10^{5}*(\frac{(M_{companion}\times M_{host})}{Separation}$ (3)

where $M_{companion}$ and $M_{host}$ are given in units of $M_{\sun}$, the
separation given in AU, and the constant chosen to give binding energy in
units of $10^{41}$ ergs. The masses of our host stars were found as described
in Section 7.3. For the companions, when no age was available, we assumed them
to be field age ($\sim$2-10 Gyr) and used the spectral type dynamical mass
bins of Dupuy & Liu (2017). For several of our systems, the host allowed for
an age range to be estimated. For these systems, we used the evolutionary
tracks of Saumon & Marley (2008) to give a range of possible masses. In order
to account for different orbital inclinations and eccentricities, we
multiplied the projected separation by 1.26 (see Fischer & Marcy 1992).
Additionally, we estimated the dissipation lifetime of each system using the
approach of Dhital et al. (2010), which estimates the length of time a system
at a given separation would survive, based on average gravitational
interactions within the galaxy. A summary of the component masses, mass
ratios, binding energies, and estimated lifetimes for all systems is shown in
Table 10.

Comparing the physical parameters of our benchmark systems with those of
stellar binary and exoplanet systems may provide clues to their possible
formation mechanisms. Figure 9 shows the separation vs. mass ratio for this
sample compared against known stellar-stellar, brown dwarf-brown dwarf, and
stellar-exoplanet systems. The exoplanet and stellar binary systems populate
two distinct regions in Figure 9, with the directly imaged exoplanets forming
a bridge between the two populations. The majority of our sample lies amongst
the stellar binary population in Figure 9, with many filling in the region
which connects with the widest and highest mass ratio directly imaged
exoplanets. Several of the systems in our sample, despite having mass ratios
more resembling directly imaged exoplanets, are found at extremely large
separations. These systems, which occupy a region of Figure 9 which is nearly
unpopulated, may be indicating the possibility for companion multiplicity,
which agrees with previous studies as discussed in Section 7.6.

Figure 10 shows the mass ratio vs. binding energy for the same sample. We can
see that the systems in this sample are beginning to fill in the sparsely
populated regions of low mass ratio and low binding energy which connect the
directly imaged exoplanets to the stellar binary population. Several systems
on Figure 10 begin to separate themselves from the other populations, with
binding energies much lower than the majority of exoplanet or stellar binary
systems.

We find that five of our hypothesized systems have binding energies below 1 x
$10^{40}$ ergs, the estimated minimum binding energy of substellar binary
systems (Burgasser et al., 2003a; Close et al., 2007), with another twenty
systems approaching this limit, with binding energies $<$5 x $10^{40}$ ergs.
These lowest binding energies may be suggestive for the need of higher
multiplicity in these systems. For example, the lowest binding energy system
of these is SDSS J134403.83+083950.9 + NLTT 35024, a T0+M5.5 pair with a
separation of $\sim$24,100 AU and a binding energy of 3.49 x $10^{39}$ ergs.
NLTT 35024 however is listed as an astrometric binary in Gaia DR3’s non-single
star solutions. Additionally, SDSS J134403.83+083950.9 was identified as a low
probability L5+T2 binary in Marocco et al. (2015), raising the possible mass
of the system even more. For more discussion on the possibility of higher
multiplicity in this sample, please refer to Section 7.4.

These low binding energies may also allow an age constraint to be placed on
the system. If the system is young, there is a possibility that gravitational
interactions have not yet had time to dissipate the pair. Indeed, several of
the lowest binding energy substellar systems known are found at young ages,
with very few found at field age (see Close et al. 2007), seeming to agree
with the picture of wide, low binding energy systems dissipating over time
through interactions with their natal environments. Previous studies have
attempted to look at ultracool companion separations as a function of age (see
Faherty et al. 2011) to investigate if there is indeed a correlation with wide
separation and youth, but found no strong correlation. The wide, low binding
energy systems in this sample are a new addition to the picture and provide an
exciting population for future analysis on the widest ultracool companions.

Further study of these low binding energy systems is required to identify any
signs of higher multiplicity in either the host or companion, as well as for
any youth diagnostics for a given system.

##  8 Conclusions

We have presented in this paper 89 new ultracool companions, including
companions to main sequence stars identified via the citizen science project
Backyard Worlds: Planet 9. We have obtained near-infrared spectra for 32 of
the companions, along with an additional 3 near-infrared spectra of ultracool
dwarfs which were originally identified as co-moving with a star, but have
since been downgraded to a low likelihood of companionship due to their low
scores with CoMover. We have also obtained both near-infrared or optical
spectra for 33 of the hosts. The spectroscopic sample of companions shows a
variety of different characteristics, including 7 young, 6 abnormally blue,
and 3 peculiarly red brown dwarfs.

Comparing our BYW benchmark sample with known companions in the literature, we
find that our sample adds 21 systems with an FGK primary, increasing the
population of ultracool companions to FGK stars by $\sim$42%. We also add 72
systems with projected separations $>$1000 AU, increasing the known population
of ultracool companions with such wide separations by 154%. The formation of
these systems remains difficult to explain, although could point towards
outward migration due to a third companion or possible dissolving moving
groups. Further study of these systems should shed light on the formation and
evolution processes of giant exoplanets and low mass stars, as they are
beginning to fill in sparsely populated spaces on separation-binding energy
diagrams.

Utilizing the properties of their host star to benchmark their properties, the
ultracool dwarfs in this BYW benchmark sample provide a wealth of information,
which can be utilized to test and constrain the atmospheric and evolutionary
models of brown dwarfs and exoplanets. Seventy-two of the ultracool dwarf
companions which are below the magnitude limit for astrometry with Gaia are
able to adopt the milliarcsecond precision parallax of their host, allowing
them to serve as parallax benchmarks needed for calibration of absolute
magnitude - spectral type trends. Thirty-seven hosts in the BYW benchmark
sample provide an age estimate for the system, breaking the mass-temperature-
age degeneracy which plagues their ultracool dwarf companion. Twenty-four
hosts in the BYW benchmark sample are FGK stars, which allow for a measure of
their elemental abundances and composition.

The BYW benchmark sample is a valuable resource for future retrieval studies,
where it is crucial to be able to compare the retrieved abundances of the
companion with the abundances of the host star. For example, this sample
provides excellent targets for comparing the C/O ratios between host and
companion, better understanding and constraining the silicate cloud properties
of ultracool atmospheres using their host’s Mg/Si ratios, or further
constraining the role metallicity plays in changing the SEDs and atmospheric
properties of ultracool dwarfs.

The authors would like to thank the numerous citizen scientist volunteers of
the Backyard Worlds: Planet 9 project for their hard work, dedication, and
contributions to this paper. The authors acknowledge support from NASA award
#80NSSC21K0402 and NASA ADAP award #80NSSC22K0491 as well as NSF award
#2009177. JF acknowledges support from the Heising Simons foundation. This
paper includes data gathered with the 6.5 meter Magellan Telescopes located at
Las Campanas Observatory, Chile. Based in part on observations obtained at the
Southern Astrophysical Research (SOAR) telescope, which is a joint project of
the Ministério da Ciência, Tecnologia e Inovações (MCTI/LNA) do Brasil, the US
National Science Foundation’s NOIRLab, the University of North Carolina at
Chapel Hill (UNC), and Michigan State University (MSU). Some of the data
presented herein were obtained at Keck Observatory, which is a private 501(c)3
non-profit organization operated as a scientific partnership among the
California Institute of Technology, the University of California, and the
National Aeronautics and Space Administration. The Observatory was made
possible by the generous financial support of the W. M. Keck Foundation. The
authors wish to recognize and acknowledge the very significant cultural role
and reverence that the summit of Maunakea has always had within the Native
Hawaiian community. We are most fortunate to have the opportunity to conduct
observations from this mountain. Some of the observations reported in this
paper were obtained with the Southern African Large Telescope (SALT). This
work presents results from the European Space Agency (ESA) space mission Gaia.
Gaia data are being processed by the Gaia Data Processing and Analysis
Consortium (DPAC). Funding for the DPAC is provided by national institutions,
in particular the institutions participating in the Gaia MultiLateral
Agreement (MLA). The Gaia mission website is https://www.cosmos.esa.int/gaia.
The Gaia archive website is https://archives.esac.esa.int/gaia. This
publication makes use of data products from the Wide-field Infrared Survey
Explorer, which is a joint project of the University of California, Los
Angeles, and the Jet Propulsion Laboratory/California Institute of Technology,
funded by the National Aeronautics and Space Administration. Several systems
in our BYW benchmark sample show spectral or kinematic indications of youth,
low metallicity, or anomalous colors. Below we enumerate on the properties of
these interesting sources and their host stars.

##  Appendix A Young Systems

###  A.1 UCAC3 113-933 + CWISE J002101.45-334631.5

The candidate host star UCAC3 113-933 is estimated to have a spectral type of
M0 (using the Gaia photometry and absolute magnitude relations of Kiman et al.
2019). UCAC3 113-933 was observed in the RAVE survey, and has an age estimate
in the literature of $\sim$155 Myr (Žerjal et al., 2017). Given the caveats of
this study, we adopt an age of 0.1–1 Gyr for UCAC3 113-933. Figure 1 shows
that UCAC3 113-933 does not appear overluminous on the Gaia CMD compared to
stars of the same color. UCAC3 113-933 has a corresponding UV source with a
near UV magnitude of NUV=21.08$\pm$0.18 mag (Bianchi et al., 2017). No X-ray,
H$\alpha$, or rotational information is available in the literature. Running
UCAC3 113-933’s full kinematics (proper motions, radial velocity, and
parallax) through BANYAN $\Sigma$, we find a field probability of 99.9$\%$.
UCAC3 113-933 has an anomalously large Gaia DR3 RUWE of 1.98, hinting at the
possibility of a hidden companion.

CWISE J002101.45-334631.5 (CW0021-3346) is a common-proper motion companion to
UCAC3 113-933. No spectrum of CW0021-3346 exists. The infrared colors of
CW0021-3346 (W1-W2=2.35$\pm$0.16 mag, $J$-$H$=0.31$\pm$0.25 mag,
$J$-W2=4.38$\pm$0.15 mag) indicate a mid to late T dwarf; while the W2
absolute magnitude (derived via adopting the parallax of UCAC3 113-933) yields
a spectral type estimate of (T1) and places the source as slightly more
luminous than typical brown dwarfs of the same color on the WISE CMD (Figure
1). While it has been shown that young L dwarfs appear fainter and redder than
their field counterparts (Faherty et al., 2012; Liu et al., 2016), the mid-
infrared absolute magnitudes of young T dwarfs are more comparable to the
older, field population (Zhang et al. 2020b). This raises the possibility that
the excess brightness of CW0021-3346 is due to binarity. Visually inspecting
the $Z$-, $Y$-, $J$-, $H$-, and $K$-band images available from the VISTA Kilo-
degree Infrared Galaxy (VIKING) survey (Edge et al., 2013), we find no
evidence of a companion, although its faint near-infrared magnitudes (J=19.9
mag) make a clear determination difficult. Additionally, the potential pair
may be too close to be resolved, leaving the possibility of binarity still
open. As has been noted in Faherty et al. (2010a), widely separated ultracool
dwarfs have a higher likelihood of being a close double themselves. Therefore
binarity is a viable option for explaining the position of CW0021-3346 in
Figure 1.

The angular separation of this system is $\sim$116″. At the distance of UCAC3
113-933 the pair have a projected physical separation of $\sim$6100 AU. With
both components in this system appearing brighter than their field
counterparts, taking into consideration UCAC3 113-933’s young age range from
Žerjal et al. (2017), we believe the UCAC3 113-933 + CW0021-3346 system to be
a 0.1–1 Gyr (M0)+(T5) widely separated binary.

###  A.2 CWISE J022454.80+152629.5 + CWISE J022454.10+152633.8

CWISE J022454.80+152629.5 (CW0224+1526A) is classified as an L6 (red) based on
the comparison of its near infrared spectrum with that of the L6 spectral
standard, shown in Figure 5. CW0224+1526A’s spectrum appears slightly redder
than the L6 spectral standard, and is indicative of a slightly more triangular
H-band, suggestive of youth. We calculated the gravity sensitive indices of
Allers & Liu (2013) for CW0224+1526A (Table 5) and find assign it a gravity
classification of INT-G. Running the kinematics of CW0224+1526A through BANYAN
$\Sigma$, and find that CW0224+1526A receives a 89.9$\%$ probability of
membership to the $\beta$ Pictoris moving group, a 10$\%$ field probability,
and a 0.1$\%$ probability of AB Doradus moving group membership.

CWISE J022454.10+152633.8 (CW0224+1526B) is a common proper motion companion
with CW0224+1526A. The infrared colors of CW0224+1526B (J-W2=4.41$\pm$0.22
mag, W1-W2=0.80$\pm$0.06 mag) are indicative of an object near the L/T
boundary, with the W1-W2 color implying a spectral type estimate of (T1). We
ran the kinematics of CW0224+1526B through BANYAN $\Sigma$ and find a 72.6$\%$
probability of membership to the $\beta$ Pictoris moving group, a field
probability of 23.1$\%$, a 2.4$\%$ probability of AB Doradus membership, and a
1.2$\%$ probability of Tucana-Horologium membership.

CW0224+1526A and CW0224+1526B are separated on the sky by an angular
separation of 11″. Using the absolute W2 - spectral type relations of
Kirkpatrick et al. (2021), we estimate CW0224+1526A to be at a distance of
$\sim$36 pc, which translates to a projected physical separation for
CW0224+1526AB of 400 AU. Both CW0224+1526 have higher probabilities of
membership to the $\beta$ Pictoris moving group, which has an age of 23$\pm$3
Myr (Mamajek & Bell, 2014). However, it would be expected for objects at the
age of $\beta$ Pictoris to have gravity classifications of VL-G, so the INT-G
classification of CW0224+1526A may be hinting that the CW0224+1526AB pair may
in fact be an interloper. With the pairs kinematics still consistant with
known young moving groups, along with the red spectral morphology and INT-G
classification of CW0224+1526A, we believe CW0224+1526AB to be a $<$1 Gyr
widely separated L6 (red)+(T1) binary.

###  A.3 HD 51400 + CWISE J065752.45+163350.2

HD 51400, is a solar-type star with a spectral type of G5 (Cannon & Pickering,
1993). On the Gaia CMD in Figure 1, HD 51400 appears normal for its spectral
type. Using the rotational period of HD 51400 from its TESS light curve, and
placing it on the color-period diagram in Figure 8, it appears to have an age
$\gtrsim$650 Myr. No information on X-ray, UV, or H$\alpha$ emission is listed
in the literature. HD 51400 has an anomalously large Gaia DR3 RUWE of 22.44,
and is listed as an astrometric binary in Gaia DR3’s Non-Single Star Solution
tables (Gaia Collaboration et al., 2023b). This strongly suggests that HD
51400 is in fact a binary itself. Running the 3D kinematics of HD 51400
through BANYAN $\Sigma$, we find a field probability of 99.9$\%$.

CWISE J065752.45+163350.2 (CW0657+1633) is a common proper motion companion to
HD 51400. We classify CW0657+1633 as spectral type L6 based on its near-
infrared spectrum (Figure 5). Calculating the gravity sensitive indices of
Allers & Liu (2013), we find that CW0657+1633 has a gravity classification of
INT-G, indicating a surface gravity lower than field age objects of the same
spectral type. On the WISE CMD (Figure 1), using the parallax of its suspected
primary, CW0657+1633 appears slightly fainter than typical objects of a
similar spectral type. Running its proper motions through BANYAN $\Sigma$, as
well as an estimated distance based on its spectral type, we find that
CW0657+1633 has a field probability of 99.9$\%$.

HD 51400 and CW0657+1633 have an angular separation of 61″. At the Gaia
distance of HD 51400, this translates to a projected physical separation of
$\sim$2,300 AU. Due to the age estimate of HD 51400 from its rotation rate, as
well as the INT-G gravity score of CW0657+1633, we believe HD 51400 +
CW0657+1633 to be a 0.65 - 1 Gyr, widely separated G5+L6 system.

###  A.4 CWISE J101533.05-111501.0AB + CWISE J101523.92$-$111539.6

CWISE J101533.05-111501.0A (1015-1115A) and CWISE J101533.05-111501.0B
(1015-1115B) form a close ($\sim$2″) binary pair. We obtained an optical
spectrum of 1015-1115A, and classify it as an M5 (Figure 4). Using the Gaia
photometry and absolute magnitude relations of Kiman et al. (2019), we
estimate 1015-1115B to have a spectral type of M5 as well. On the Gaia CMD
(Figure 1), we find that while 1015-1115B appears normal for its color,
1015-1115A appears to be overluminous. In addition, the spectrum of 1015-1115A
in Figure 4 shows noticeable H$\alpha$ emission. There is a near UV source
associated with the binary pair (NUV=22.51$\pm$0.42 mag; Bianchi et al. 2017),
although there is no associated X-ray emission. We find a rotational period of
Prot=0.35$\pm$0.11 days from the TESS light curve, however due to the large
pixel size of TESS both 1015-115A and 1015-1115B appear in the same pixel.
Making the interpretation of this rotation period more complicated is that
1015-1115A has a large Gaia DR3 RUWE value of 1.67, indicating the possibility
of a hidden companion. On the color-period diagram shown in Figure 8,
1015-1115A appears to be consistent with an age of only a few hundred of
millions of years. We ran 1015-1115A through BANYAN $\Sigma$ (as 1015-1115A
was the only star in the pair with full kinematics in Gaia) and find that it
receives a probability of membership to the TW Hydrae Association (TWA) of
85.2$\%$ and a field probability of 14.8$\%$ .

CWISE J101523.92$-$111539.6 (CW1015-1115) is a common proper motion companion
with the 1015-1115A and 1015-1115B pair. We assign CW1015$-$1115 a spectral
type of L5 based on its near infrared spectrum in Figure 5. It can be seen
that the spectrum of CW1015$-$1115 appears slightly redder compared to the L5
standard in Figure 5, with an excess of flux in the $H$\- and $K$-bands. The
peak of the $H$-band of CW1015$-$1115’s spectrum also displays a distinctly
more triangular shape. Using the gravity sensitive indices of Allers & Liu
(2013) (listed in Table 5), we find that CW1015$-$1115 receives a gravity
score of INT-G, indicating a lower surface gravity than the field population.
CW1015$-$1115’s location on the WISE CMD in Figure 1 is in good agreement with
other mid L dwarfs, although it appears slightly brighter than the field L5s.
Young L dwarfs appear fainter in the near-infrared than their field
counterparts, though have been shown to be overluminous in mid-infrared bands
(Liu et al., 2016), so CW1015-1115’s elevated position also hints at a
possible younger than field age. This overluminosity may instead be indicative
of CW1015-1115 being a close binary itself, however, which is thought to be
more common in widely separated systems such as this (Faherty et al., 2010a).
Running the kinematics of CW1015$-$1115 through BANYAN $\Sigma$, using an
estimated parallax of 29$\pm$6 mas (derived from spectral type - MW2
relations), we find a membership probability of 86.1$\%$ for TWA and a field
probability of 13.9$\%$, similar to the results for 1015-1115A. Running BANYAN
$\Sigma$ on CW1015-1115 without an estimated parallax results in a 71.2$\%$
probability of TWA membership, a 1.5$\%$ probability of membership to the
Carina Near Moving Group, and a 27.3$\%$ field probability.

CW1015$-$1115 has an angular separation from 1015-1115A and 1015-1115B of
$\sim$140″. At the 1015-1115A, this translates to a projected physical
separation of $\sim$5,800 AU. While the system receives a high membership
probability of TWA membership, the INT-G gravity score of CW1015$-$1115 is at
odds with this assessment. Faherty et al. (2016) found that all TWA members of
spectral type M7 or later received gravity scores of VL-G. While not
impossible, as other physical processes may be impacting the relevant spectral
features, it is unlikely for CW1015-1115 to receive an INT-G gravity score at
the age of TWA, and may be an interloper to the YMG. Future follow up is
required to further assess the spectral and kinematic indications of youth for
CW1015-1115. However, due to 1015-1115A’s H$\alpha$ emission and fast rotation
rate, the redder color, triangular H band, and INT-G gravity score of
CW1015$-$1115’s spectrum, we believe the 1015-1115A + 1015-1115B +
CW1015$-$1115 to be a 0.1-1 Gyr, widely separated M5+(M5)+L5$\beta$ triple.

###  A.5 LSPM J1417+0418 + CWISE J141737.21+041847.2

LSPM J1417+0418 is a known high proper motion star. While both the Gaia
photometry - spectral type relations of Kiman et al. (2019), as well as the
Gaia CMD location in Figure 1, place LSPM J1417+0418 as a mid-M dwarf, visual
inspection using the WiseView tool leads us to believe this to be a result of
contamination. WiseView includes a feature which allows for an overlay to be
placed on the image, representing the parallax and proper motion of Gaia
sources within the field of view. In looking at the location of LSPM
J1417+0418 using this feature, there does not appear to be a moving WISE
source associated with it, only a stationary source. This lack of an
associated infrared source to a Gaia source is sometimes seen with white
dwarfs, which are not particularly bright in these wavelengths. Inspecting
higher resolution images from DSS, Pan-STARRS, and UKIDSS, there do indeed
appear to be two different sources: one which is believed to be a background
galaxy, SDSS J141737.75+041823.9 (Vasconcellos et al., 2011), and one with a
similar proper motion as the Gaia source. While the extended source nature of
the background object potentially contaminates the photometry of LSPM
J1417+0418, making any spectral type estimating difficult, we believe LSPM
J1417+0418 to be a possible white dwarf based on the lack of an associated
moving infrared source. Further observations of LSPM J1417+0418 are required
in order to confirm. As the Gaia RUWE for LSPM J1417+0418 of RUWE=1.29 falls
below the nominal 1.4 limit for a ’well behaved’ single star (see Stassun &
Torres 2021), the Gaia astrometry for LSPM J1417+0418 is likely unaffected by
the contamination. Running the proper motions and parallax for LSPM J1417+0418
through BANYAN $\Sigma$, we find a 99.9% field probability.

CWISE J141737.21+041847.2 (CW1417+0418) is a common proper motion companion
with LSPM J1417+0418. CW1417+0418 is found to also have a corresponding Gaia
source, whose parallax of $\varpi$=16.01$\pm$0.45 mas is similar to the
parallax of LSPM J1417+0418 of $\varpi$=17.07$\pm$0.27. The near infrared
spectrum of CW1417+0418 is best matched by the near infrared spectrum of the
M8 standard (Figure 5). Using the indices of Allers & Liu (2013) we find that
CW1417+0418 receives a gravity score of INT-G (listed in Table 5), indicating
an object with a surface gravity lower than the field population. Looking at
CW1417+0418 on the WISE CMD in Figure 1, we see that it appears near normal
for its spectral type. On the Gaia CMD in Figure 1 however, CW1417+0418
appears slightly redder than objects of the same spectral type, placing it
just to the right of the late M main sequence. Running the Gaia proper motions
and parallax of CW1417+0418 through BANYAN $\Sigma$, we find a 99.9% field
probability.

LSPM J1417+0418 and CW1417+0418 are separated on the sky by an angular
separation of $\sim$24″. At the Gaia distance of LSPM J1417+0418, this
translates to a projected physical separation of $\sim$1,400 AU. While the
gravity sensitive indices for CW1417+0418 are suggestive of an object with
intermediate gravity, the unknown nature and age of its host star makes an
accurate determination difficult. It is unlikely that LSPM J1417+0418 is
young, however, if its white dwarf nature is confirmed. This hints that
CW1417+0418’s INT-G classification is not due to youth, but to other processes
subtly affecting the observed spectrum as has been found for several other M
dwarfs (c.f. TRAPPIST-1 Gonzales et al. 2019) We therefore believe LSPM
J1417+0418 + CW1417+0418 to be a widely separated (WD)+M8$\beta$.

###  A.6 UCAC4 840-013771 + CWISE J162511.27+774946.8

UCAC4 840-013771 is a high proper motion star which we classify as spectral
type M2 based on its optical spectrum, obtained from the CFHT archive (Figure
4). While UCAC4 840-013771’s spectrum matches the M2 spectral standard between
$\sim$6000 Å and $\sim$7600 Å in Figure 4, it is heavily suppressed compared
to the spectral standard $>$7600 Å. We note this suppression may be due to
reduction or observation problems. Figure 4 also shows the spectrum of UCAC4
840-013771 exhibits strong H$\alpha$ emission. Riaz et al. (2006) report an
H$\alpha$ equivalent width of 4.5 Å, agreeing with the strong H$\alpha$
emission in the observed spectrum. UCAC4 840-013771 is an X-ray and UV
emitter, with an observed X-ray flux of FX= 6.72 $\times$ 10-13 erg s-1 cm-1
(Freund et al., 2022) and UV magnitudes of FUV=21.04$\pm$0.24 mag and
NUV=19.99$\pm$0.09 mag (Bianchi et al., 2017). Using the TESS lightcurve of
UCAC4 840-013771, we find a rotation period of Prot=1.3 days. On the color-
period diagram shown in Figure 8, UCAC4 840-013771 appears to be consistent
with an age of at most a few hundred million years. The light curve of UCAC4
840-013771 also shows evidence of substantial flare activity. Looking at its
position on the Gaia CMD in Figure 1, UCAC4 840-013771 appears overluminous
compared to other stars of the same color. The Gaia DR3 RUWE for UCAC4
840-013771 is larger than typical for single stars (RUWE=1.61), which could
possibly indicate a hidden companion, or be a result of variability due to the
bright flares evident in its TESS light curve. Running UCAC4 840-013771
through BANYAN $\Sigma$ with the full kinematics from Gaia DR3, we get a
99.9$\%$ field probability.

CWISE J162511.27+774946.8 (CW1625+7749) is a common proper motion companion
with UCAC4 840-013771. We classify CW1625+7749 as spectral type L4 from its
near infrared spectrum (Figure 5). The Keck/NIRES spectrum for this source was
taken through variable and thick clouds, and its high declination resulted in
a large airmass difference between it ($z$ = 1.9) and the nearest A0 V
telluric calibrator ($z$ = 1.5). As such, both the relative calibration across
orders (e.g., the anomalously bright K-band) and the correction of telluric
absorption (in particular the 1.1–1.5 $\mu$m band) are inaccurate. In order to
properly spectral type CW1625+7749, we followed the procedure used in the past
on extremely red L dwarfs (see for example Marocco et al. 2014; Cruz et al.
2018): each of the 3 bands were separated and individually normalized, then
compared with the spectral standards treated in the same manner. This band-by-
band classification results in a classification of L2. Calculating the indices
of Allers & Liu (2013) (Table 5), we find that CW1625+7749 receives a gravity
score of FLD-G, indicating an object of field age surface gravity. We note
though that this score may not be accurate, due to potential reduction issues
caused by strong telluric absorption and flux calibration issues. Looking at
the position of CW1625+7749 on the WISE CMD, it does not appear to be under or
over luminous compared with field objects of the same spectral type. We ran
CW1625+7749 through BANYAN $\Sigma$, using an estimated parallax of
15.20$\pm$3.04 mas derived from spectral type - MW2 empirical relations, and
find a 99.9$\%$ field probability.

CW1625+7749 and UCAC4 840-013771 are separated on the sky by an angular
distance of $\sim$313″. At the Gaia DR3 distance of UCAC4 840-013771, this
translates to a projected physical separation of $\sim$19,000 AU. Given UCAC4
840-013771’s elevated CMD, H$\alpha$, X-ray and UV emission, as well as its
fast rotation rate and flares, we believe UCAC4 840-013771 + CW1625+7749 to be
a 0.1–1 Gyr M2+L2 wide binary.

###  A.7 TYC 5213-545-1 + CWISE J214129.80-024623.6

TYC 5213-545-1 is classified as a K0 dwarf (Xiang et al., 2019). Using the
TESS light curve of TYC 5213-545-1, we find a rotational period of Prot=14.3
days, with hints of multiperiodic oscillations. No H$\alpha$, X-ray, or UV
data is reported in the literature. Žerjal et al. (2017) report an age for TYC
5213-545-1 of $\sim$550 Myr based on the Ca II infrared triplet equivalent
width excess, and we adopt an age range of 0.1–1 Gyr. On the Gaia CMD, TYC
5213-545-1 does not appear overluminous compared to stars of similar color.
Its Gaia DR3 RUWE is near one (RUWE=0.95). Running the full Gaia kinematics of
TYC 5213-545-1 through BANYAN $\Sigma$, we find a 99.9$\%$ field probability.

CWISE J214129.80-024623.6 (CW2141-0246) is a common proper motion companion
with TYC 5213-545-1. We assign CW2141-0246 a spectral type of L4 based on its
near infrared spectrum. In Figure 5 the spectrum of CW2141-0246 appears more
red than the standard L4, with its $H$ and $K$ bands slightly enhanced. The
peak of CW2141-0246’s $H$-band does not show the typical triangular shape, but
is instead much flatter than the spectral standard in Figure 5. CW2141-0246
receives a gravity classification of FLD-G using the indices of Allers & Liu
(2013), indicating an object with field age surface gravity. The WISE CMD in
Figure 1 shows that CW2141-0246 is in good agreement with typical field dwarfs
of the same color. Table 6 shows CW2141-0246 is consistent with being a weak
spectral binary candidate. BANYAN $\Sigma$, with an estimated parallax of
15.9$\pm$3.2 mas derived from empirical spectral type - MW2 relations, we find
a field probability of 99.9$\%$.

TYC 5213-545-1 and CW2141-0246 are separated on the sky by $\sim$114″. At the
Gaia DR3 distance of TYC 5213-545-1, this translates to a projected physical
separation of $\sim$9,000 AU. Given the red spectrum of CW2141-0246, along
with the age range of 0.1–1 Gyr from Žerjal et al. (2017), we believe TYC
5213-545-1 + CW2141-0246 to be a 0.1–1 Gyr K0+L4 wide binary.

###  A.8 GJ 900 + CWISE J233531.55+014219.6

GJ 900 is a multiple system, where the components A, B, and C have spectral
types K7, M4, and M6 respectively (Malogolovets et al., 2007). GJ 900 is a
hierarchical triple: GJ 900A and GJ 900BC orbiting a common center of mass
with a period of $\sim$80 years, and GJ 900BC form a close binary pair with an
orbital period of $\sim$20 yrs . Malogolovets et al. (2007) also report the
detection of two fainter objects in 2MASS images which are theorized to
potentially be late M dwarfs comprising components D and E of the GJ900
system. No follow-up observations have confirmed these sources as part of the
GJ 900 system, and are likely not associated. On the Gaia CMD in Figure 1, GJ
900 appears to be overluminous, as is expected for multiple systems. Gagné et
al. (2018) assign GJ 900 a bona-fide membership of the Carina-Near moving
group (200$\pm$50 Myr). We reran GJ 900 through BANYAN $\Sigma$ using the
updated kinematics from Gaia DR3 and find a 99.7$\%$ membership probability to
Carina-Near and a 0.3$\%$ field probability, which is in good agreement with
Gagné et al. (2018). Using its light curve from TESS, we find a rotational
period of 11.92 days, along with a smaller amplitude periodicity of $\sim$8
days. The light curve of GJ 900 also shows evidence of flares. On the color-
period diagram in Figure 8, GJ 900’s position is consistent with its given age
of 200$\pm$50 Myr. The GJ 900 system is a source of X-ray emission, with an
observed flux of 9.13 $\times$ $10^{2}$ $mW/m^{2}$ (Salvato et al., 2018), as
well as a source of UV emission, with observed magnitudes of
FUV=19.94$\pm$0.14 mag and NUV=17.52$\pm$0.03 mag (Bianchi et al., 2017). GJ
900 has multiple reported metallicity measurements in the literature ranging
from [Fe/H] of –0.6 to +0.3 (Alonso et al., 1996; Ammons et al., 2006;
Houdebine et al., 2016; Soubiran et al., 2016; Houdebine et al., 2019; Gaia
Collaboration et al., 2023a).

CWISE J233531.55+014219.6 (CW2335+0142) is a common proper motion companion
with GJ 900. The near infrared spectrum of CW2335+0142 in Figure 5, while low
signal-to-noise due to observational challenges, is consistent with a cold
object of late-T spectral type, and is best fit as spectral type T9. On the
WISE CMD in Figure 1, CW2335+0142 does not appear to be over- or under-
luminous when compared to objects of a similar color. Running the kinematics
and estimated photometric distance CW2335+0142 through BANYAN $\Sigma$, we
find a 73.2$\%$ probability of $\beta$ Pictoris moving group membership, a
12.7$\%$ probability of AB Doradus Moving Group membership, 6.4$\%$
probability of Argus Moving Group membership, and a 7.6$\%$ field probability.
While these membership probabilities for CW2335+0142 are not in agreement with
those of its suspected host (99.7$\%$ Carina-near), this may be due to the
large proper motion uncertainties ($\sim$85 mas/yr) used for CW2335+0142 in
BANYAN $\Sigma$. However, with only a 7.6$\%$ field probability, the
kinematics of CW2335+0142 are still strongly suggestive of it belonging to a
young population. Further observations of CW2335+0142, including higher
accuracy proper motions, are required in order to confirm any potential
memberships.

GJ 900 and CW2335+0142 are separated on the sky by an angular separation of
$\sim$587″. At the distance of GJ 900, this translates to a projected physical
separation of $\sim$12,000 AU. Despite the disagreement of YMG memberships
from BANYAN $\Sigma$, the GJ 900 + CW2335+0142 pair still receive a 99.9$\%$
probability of being a co-moving pair via the CoMover code. Given the high co-
moving probability from CoMover, along with GJ 900’s membership of the Carina-
Near moving group, we believe GJ 900 + CW2335+0142 to be a 200$\pm$50 Myr
widely-separated (K7+M4+M6)+T9 multiple system.

One of the only comparably cold, widely separated substellar companions known
is COCONUTS-2b, a T9 with and age of 150-800 Myr and a mass of
6.3${}^{+1.5}_{-1.9}$ $M_{\rm Jup}$(Zhang et al., 2021b). If confirmed to
belong to the Carina-Near moving group, CW2335+0142 would also be a planetary
mass object, adding to the small number of widely separated, cold exoplanets.

##  Appendix B Unusually Blue Spectra

The brown dwarf population has several subpopulations that showcase how
varying secondary parameters might alter observable features. One such
subpopulation is subdwarfs – low metallicity objects which present enhanced
collision induced H2 absorption resulting in a bluer than normal near-infrared
spectrum (Linsky, 1969; Saumon et al., 1994). In addition to low metallicity,
in L dwarfs, abnormally blue near-infrared colors have been shown to be a
possible indication of unresolved binarity with a T dwarf component
(Burgasser, 2007; Bardalez Gagliuffi et al., 2014). There are also some
ultracool dwarfs which present bluer than normal spectra, but do not show
spectral indications of low metallicity or binarity (Gizis et al., 2000; Cruz
et al., 2003; Kirkpatrick et al., 2010b). It has been hypothesized that large
grained or thin condensate clouds may be responsible for these blue L-dwarfs
(Burgasser et al., 2008). Below we discuss the systems which presented bluer
than expected near-infrared spectra.

###  B.1 G 73-59 + CWISE J022737.75+083008.8

G 73-59 is a high proper motion star, which our Kast optical spectrum
indicates a spectral type of M3. On the Gaia CMD in Figure 1, G 73-59 appears
to have the same absolute magnitude as stars with a similar color. The Gaia
DR3 RUWE of G 73-59 is abnormally large (RUWE=15.53), possibly indicating a
hidden binary companion. Using the Gaia position, proper motion, radial
velocity, and parallax of G 73-59, we calculated the 3D galactic velocity
($U_{lsr}$, $V_{lsr}$, $W_{lsr}$), correcting for solar motion using the
values of $U_{\sun},V_{\sun},W_{\sun}$=(10.0, 5.25, 7.17) km s-1 from Dehnen &
Binney (1998). The resultant velocities for G 73-59 are $U_{lsr}$, $V_{lsr}$,
$W_{lsr}$=(-66.8, -26.5, -39.1) km s-1 placing it as kinematically similar to
stars found within the thick disk (Bensby et al., 2003). No metallicity
measurements currently exist in the literature for G 73-59.

CWISE J022737.75+083008.8 (CW0227+0830) is a common proper motion companion
with G 73-59. The near infrared spectrum of CW0227+0830 in Figure 5 classifies
it as having a spectral type of L5 (blue). In Figure 5, the spectrum of
CW0227+0830 is noticeably bluer than typical field dwarfs of the same spectral
type. CW0227+0830 has a $J$-$K$ color of 1.29$\pm$0.24 mag, which when
compared to the average L5 color of $J$-$K$=1.75$\pm$0.22 mag from Faherty et
al. (2016), we see that CW0227+0830 is a blue color outlier, agreeing with its
observed spectrum. Looking at CW0227+0830 on the WISE CMD in Figure 5,
although its location is consistent with typical field L dwarfs, it appears
slightly fainter, placing it closer to later-type mid-L dwarfs. The tangential
velocity of CW0227+0830, $v_{tan}$=63$\pm$2 km s-1, is much more similar to
other blue outliers than to typical field dwarfs, which have average
tangential velocities of $v_{tan}$=53$\pm$47 km s-1 and $v_{tan}$=26$\pm$19 km
s-1, respectively (Faherty et al., 2009), hinting that CW0227+0830 may belong
to an older population. We also tested the possibility for the bluer spectral
morphology of CW0227+0830 to be due to an unresolved companion, as described
in Section 5.5. Table 6 shows that CW0227+0830 meets the criteria for a weak
spectral binary candidate.

G 73-59 and CW0227+0830 are separated on the sky by an angular separation of
$\sim$16″. At the distance of G 73-59, that translates to a projected physical
separation of $\sim$550 AU. Given the high kinematics of G 73-59, as well as
the blue spectral morphology of CW0227+0830, we believe G 73-59 + CW0227+0830
to be an older widely separated M3+L5 (blue) pair.

###  B.2 LAMOST J0626+5933AB + CWISE J062648.96+594129.2

LAMOST J062631.15+593341.3 and LAMOST J062631.34+593343.6 (hereafter referred
to as LAMOST J0626+5933A and LAMOST J0626+5933B) form a close binary pair,
with an angular separation of 2.8″ and spectral types M1 and M2 respectively
(Xiang et al., 2019). Only LAMOST J0626+5933B has a corresponding entry in
Gaia, so looking at only LAMOST J0626+5933B on the Gaia CMD in Figure 5, we
see that it appears to be slightly fainter than stars of a similar color.
Xiang et al. (2019) report multiple metallicities for both components, but
only [Fe/H]=-0.88$\pm$0.37 for LAMOST J0626+5933B is listed as reliable. Ding
et al. (2022) report metallicities for LAMOST J0626+5933A and LAMOST
J0626+5933B (with repeated measurements) of [Fe/H]=-0.45$\pm$0.26 and
[Fe/H]=-0.43$\pm$0.26 respectively. Additionally, LAMOST J0626+5933B has a
global metallicity measurement of [M/H]=$-0.81^{+0.14}_{-0.20}$ from the Gaia
RVS spectrum, as well as an $\alpha$ abundance of
[$\alpha$/Fe]=$0.29^{+0.07}_{-0.06}$. There are no reported UV or X-ray
detections in the literature, and no TESS light curve is available. We
calculated the 3D velocity of LAMOST J0626+5933B as described in Section B.1,
and obtain the values $U_{lsr}$, $V_{lsr}$, $W_{lsr}$=(-85.2, 6.9, -10.4) km
s-1 placing it as kinematically similar to stars within the thick disk (Bensby
et al., 2003).

CWISE J062648.96+594129.2 (CW0626+5941) is a common proper motion companion
with LAMOST J0626+5933AB. The near infrared spectrum of CW0626+5941, shown in
Figure 5, classifies it as a spectral type of L2 (blue), as its spectrum is
slightly bluer than typical field dwarfs of the same spectral type.
CW0626+5941 has a $J$-$K$ color of 1.45$\pm$0.19 mag, slightly bluer than the
average L2 color of $J$-$K$=1.51$\pm$0.21 mag (Faherty et al., 2016), however
it is not blue enough to classify it as a photometric outlier. The tangential
velocity of CW0626+59441, $v_{tan}$=66$\pm$4 km s-1, is kinematically similar
to the average $v_{tan}$=53$\pm$47 km s-1 of known blue outliers (Faherty et
al., 2009).

LAMOST J0626+5933AB and CW0626+5941 are separated on the sky by an angular
separation of $\sim$488″. At the Gaia distance of LAMOST J0626+5933B, this
translates to a projected physical separation of $\sim$27,000 AU. Given the
reported sub-solar metallicities and large 3D velocity for LAMOST
J0626+5933AB, as well as the slightly blue spectral morphology of CW0626+5941
and its $v_{tan}$ similar to blue outliers, we believe LAMOST J0626+5933AB +
CW0626+5941 to be a widely separated M1+M2+L2 (blue) triple.

###  B.3 WISE J062727.34$-$002826.8 + CWISE J062725.95$-$002843.8

CWISE J062727.34$-$002826.8 (CW0627-0028A) is classified as a T0 (blue) dwarf
via comparison of its near infrared spectrum with the near infrared spectral
standards, shown in Figure 5. The spectrum of CW0627$-$0028A in Figure 5 is
much bluer, however, than typical field dwarfs of the same spectral type. The
$J$-$K$ color of CW0627$-$0028A (1.02$\pm$0.32 mag) is much bluer than the
average T0 color of $J$-$K$=1.63$\pm$0.40 mag (Faherty et al., 2009), agreeing
with the bluer spectral morphology of CW0627$-$0028A. The mid infrared color
of CW0627$-$0028A, W1-W2=0.36$\pm$0.04 mag, is also bluer than expected, being
more consistent with mid-L dwarfs. We investigated the possibility that the
bluer colors of CW0627-0028A are due to an unresolved companion, as described
in Section 5.5. This analysis (Table 6) shows that CW0627-0028A does appear
consistent with being a weak spectral binary candidate. Using the best fit
spectral type for CW0627$-$0028A of T0, along with a distance estimate from
the absolute magnitude - spectral type relations of Kirkpatrick et al. (2021),
we estimate CW0627$-$0028A to be at distance of 35$\pm$7 pc.

CWISE J062725.95$-$002843.8 (CW0627$-$0028B) is a common proper motion
companion with CW0627-0028A. The mid-infrared color of CW0627$-$0028B,
W1-W2=0.34$\pm$0.05 mag, is indicative of a mid-L spectral type. However,
assuming CW062$-$0028A and CW0627$-$0028B are coeval, CW0627$-$0028B likely
also has bluer colors for its spectral type than typical field dwarfs, making
an accurate spectral type estimate difficult. As the mid- and near-infrared
colors of CW0627$-$0028A and CW0627$-$0028B are similar, we estimate
CW0627$-$0028B to likely have a spectral type of early-T. Further
spectroscopic observations are required to confirm the spectral type of
CW0627$-$0028B, and to determine if it displays a bluer spectrum similar to
CW0627$-$0028A.

CW0627$-$0028A and CW0627$-$0028B are separated on the sky by an angular
separation of $\sim$27″. While no parallax for either object exists, using the
estimated distance of CW0627$-$0028A, this translates to a projected physical
separation of $\sim$860 AU. This separation would make the CW0627$-$0028AB
system the widest known T+T binary system, exceeding ULAS J020529.62+142114.0,
a T1+T3 binary with a projected separation of 71 AU (Day-Jones et al., 2013).
Additionally, the CW0627$-$0028AB system would be the widest known field age
substellar binary, which is currently CWISE J0146-0508AB (L4+L8 (blue)
spectral types) at a separation of $\sim$129 AU (Softich et al., 2022).
However, an accurate parallax for the system is required. Given the blue near
infrared spectrum of CW0627$-$0028A, as well as the high co-moving probability
of the pair, we believe CW0627$-$0028AB to be a widely separated T0 (blue)+(T)
binary.

###  B.4 LP 677-81 + CWISE J132857.58$-$063747.4

LP 677-81 is an estimated M3 dwarf using the Gaia photometric relations of
Kiman et al. (2019). Looking at LP 677-81 on the Gaia CMD in Figure 1, we see
that it is slightly fainter compared to other stars of the same color. No TESS
light curve is available for LP 677-81. However, Reinhold & Hekker (2020a)
report a rotational period of $P_{rot}=23\pm 23$ days. From looking at LP
677-81 on the color-period diagram shown in Figure 8, we find that its age
appears to be at least $>$650 Myr. No metallicity measurements have been
reported in the literature for LP 677-81, and its Gaia DR3 RUWE of 1.07
indicates that it is likely not an unresolved binary. Using the Gaia
kinematics of LP 677-81, we calculated its 3D space motion as described in
Section B.1, finding the velocities $U_{lsr}$, $V_{lsr}$, $W_{lsr}$=(62.9,
-51.1, 0.7) km s-1 making LP 677-81 kinematically similar to the thick disk
population (Bensby et al., 2003). .

CWISE J132857.58$-$063747.4 (CW1328$-$0637) is a common proper motion
companion with LP 677-81. We classify CW1328$-$0637 as spectral type L3 (blue)
based on its near infrared spectrum in Figure 5. As can be seen in Figure 5,
the spectrum of CW1328$-$0637 is much bluer than typical field dwarfs of the
same spectral type. The near infrared color of CW1328$-$0627,
$J$-$K$=0.96$\pm$0.07 mag, indicates it as a blue color outlier when compared
with the average for that spectral type, $J$-$K$=1.61$\pm$0.22 mag (Faherty et
al., 2016), which is in agreement with the blue morphology of its observed
spectrum. Looking at CW1328$-$0637 in the WISE CMD in Figure 1 and adopting
the parallax of its suspected host, we see that it appears faint for its
spectral type, with an absolute magnitude placing it among mid to late L
dwarfs. We find that CW1328-0637 is a weak spectral binary candidate (Table
6), however Bardalez Gagliuffi et al. (2014) found that blue L-dwarfs are a
major contaminant when investigating spectral binaries. Adopting the Gaia
parallax of LP 677-81 we find that CW1328$-$0637 has a tangential velocity of
$v_{tan}=77\pm 12$ km s-1, similar to the tangential velocity of other blue
color outliers (Faherty et al., 2009).

LP 677-81 and CW1328$-$0637 are separated on the sky by an angular separation
of $\sim$19″. At the distance of LP 677-81, this translates to a projected
physical separation of $\sim$1400 AU. Due to LP 677-81 appearing fainter in
the Gaia CMD, along with the bluer near infrared spectrum of CW1328$-$0637, we
believe LP 677-81 + CW1328$-$0637 to be a (M3)+L3 (blue) widely separated
binary.

###  B.5 CWISE J160653.16$-$103210.6 + CWISE J160654.19$-$103214.7

CWISE J160653.16$-$103210.6 (CW1606$-$1032A) is a high proper motion star with
an optical classification of M9 based on our Kast optical spectrum. Looking at
CW1606$-$1032A on the Gaia CMD in Figure 1, we see that it is in good
agreement with typical field stars of the same color. No information on
rotational period, metallicity, H$\alpha$, X-ray, or UV emission is available
in the literature. While we are not able to compute the full 3D velocity of
CW1606$-$1032A, as no radial velocity measurement exists, using the Gaia
distance for CW1606$-$1032A we find that it has a tangential velocity of
$v_{tan}$=25 km s-1. Comparing this tangential velocity with the average
tangential velocities computed in Faherty et al. (2009), we see that
CW1606$-$1032A has a tangential velocity more similar to field late M dwarfs.
The near infrared color of CW1606$-$1032A, $J$-$K$=1.03$\pm$0.06 mag, is also
normal for typical field objects of the same spectral type.

CWISE J160654.19$-$103214.7 (CW1606$-$1032B) is a common proper motion
companion with CW1606$-$1032A. Comparing the near infrared spectrum of
CW1606$-$1032B in Figure 5 with the near infrared spectral standards, we
classify CW1606$-$1032B as spectral type L7 (blue), as the spectrum of
CW1606-1032B is much bluer than typical field dwarfs (Figure 5). Additionally,
the spectrum of CW1606$-$1032B is suggestive of stronger FeH absorption bands,
a feature observed in objects with sub-solar metallicity (Burgasser et al.,
2003b; Kirkpatrick et al., 2010b). Looking at CW1606$-$1032B on the WISE CMD
using the parallax of its suspected host, we see that it lies amongst other
late-type L dwarfs. The near infrared color of CW1606$-$1032B,
$J$-$K$=1.58$\pm$0.05 mag, while bluer than average field objects of the same
spectral type, is not enough to classify it as a blue outlier. We find that
CW1606$-$1032B is a weak spectral binary candidate (Table 6) following the
methods described in Section 5.5, however as noted in Bardalez Gagliuffi et
al. (2014), blue L-dwarfs are a typical spectral binary contaminant.

CW1606$-$1032A and CW1606$-$1032B are separated on the sky by an angular
separation of $\sim$16″. At the Gaia distance of CW1606$-$1032A, this
translates into a projected physical separation of $\sim$640 AU. While the
near infrared spectrum of CW1606$-$1032B resembles the morphology observed in
some sub-solar metallicity objects, the kinematics and colors of
CW1606$-$1032A are typical for field age dwarfs. We therefore suggest that
CW1606$-$1032AB are a field-age M9+L7 (blue) widely separated binary.

###  B.6 SRC J2029$-$7910 + CWISE J202934.80$-$791013.1

SRC J2029$-$7910 is a high proper motion star which we classify as spectral
type M6 (blue) based on its near infrared spectrum. The spectrum of SCR
J2029$-$7910 in Figure 3 however has a much bluer appearance than the spectral
standard. Looking at SCR J2029$-$7910 on the Gaia CMD in Figure 1, we see that
it appears fainter compared to objects of the same spectral type. The Gaia
RUWE of SCR J2029$-$7910 is anomalously high (RUWE=1.88), hinting at the
possible existence of a hidden companion. Additionally, SCR J2029$-$7910 has
an entry in the Gaia DR3 non-single stars table, which flags SCR J2029$-$7910
as a likely binary through the analysis of its astrometric acceleration. Using
the Gaia DR3 MSC model (which treats sources as though they were an unresolved
binary), SCR J2029$-$1032 has a sub-solar metallicity of
[M/H]$\approx\\-0.20$. No information on rotation period, X-Ray, or UV
emission is available in the literature. Using the full Gaia kinematics for
SCR J2029$-$7910, we calculated its 3D space motions as described in Section
B.1, obtaining the velocities $U_{lsr}$, $V_{lsr}$, $W_{lsr}$=(-34.8, -37.6,
-29.2) km s-1. This 3D velocity places SCR J2029-7910 as more kinematically
consistent with the thin disk population (Bensby et al., 2003).

CWISE J202934.80$-$791013.1 (CW2029$-$7910) is a common proper motion
companion with SCR J2029$-$7910\. CW2029$-$7910 was also identified as being
statistically the same distance as SCR J2029$-$7910, with parallaxes of
$\varpi$=19.63$\pm$0.66 mas and $\varpi$=20.12$\pm$0.04 mas, respectively.
Reylé (2018a) assigned CW2029-7910 a phototype of L0.5 young. However we
classify CW2029$-$7910 as spectral type L1 (blue) based on its near infrared
spectrum in Figure 5. Compared to the near infrared spectral standard as shown
in Figure 5, the spectrum of CW2029$-$7910 is much bluer, and is suggestive of
stronger FeH bands and K I lines, signatures of lower metallicity. Looking at
CW2029$-$7910 on the WISE CMD in Figure 1, we find that it appears slightly
faint for its spectral type. On the Gaia CMD in Figure 1 however,
CW2029$-$7910 appears near normal compared to field objects of the same color.
CW2029-7910 receives gravity classification of FLD-G (Table 5) from the
indices of Allers & Liu (2013). No metallicity, rotational period, X-Ray, or
UV emission data are available in the literature. CW2029$-$7910 has a near
infrared color of $J$-$K$=1.02$\pm$0.17 mag, placing it as a borderline blue
color outlier. Table 6 shows that CW2029$-$7910 is a strong spectral binary
candidate. The tangential velocity of CW2029-7910 is $v_{tan}$=71.5 km s-1,
similar to other color outliers in (Faherty et al., 2009), suggestive that it
may be an older dwarf star.

SCR J2029$-$7910 and CW2029$-$7910 are separated on the sky by an angular
separation of $\sim$34″. At the Gaia distance of SCR J2029$-$7910, this
translates to a projected physical separation of $\sim$1700 AU. Given the
bluer spectral morphologies of SCR J2029$-$7910 and CW2029$-$7910, as well as
the sub-solar metallicity of SCR J2029$-$7910, we believe that SCR
J2029$-$7910 + CW2029$-$7910 are a widely separated low-metallicity M6
(blue)+L1 (blue) system.

##  Appendix C Unusually Red Spectra

There exist some brown dwarfs that display red near-infrared colors, yet show
no indications of youth. It has been suggested that an excess of dust,
consisting of micron sized grains, high in the brown dwarf atmospheres would
preferentially suppress flux at shorter wavelengths, resulting in the observed
reddening (Cushing et al., 2008; Marocco et al., 2014; Hiranaka et al., 2016).
Viewing angle has also been shown to be correlated with color anomalies, where
objects viewed equator-on have been observed to be redder than other objects
viewed at different latitudes (Vos et al., 2017). This correlation has been
explained by cloud properties, as equatorial latitudes are cloudier than polar
latitudes (Suárez et al., 2023). Below we discuss the objects in our sample
which have redder than expected near-infrared spectra which cannot be
attributed to youth.

###  C.1 LSPM J0738+5254 + CWISE J073831.31+525453.7

LSPM J0738+5254 is classified as spectral type M3 (Luo et al., 2019). On the
Gaia CMD in Figure 1, LSPM J0738+5254 appears normal for its spectral type.
Ding et al. (2022) find a near solar metallicity of [Fe/H]=$0.12\pm 0.26$,
whereas Zhang et al. (2020a) find a metallicity of [M/H]=$-0.62\pm 0.13$. Lu
et al. (2019) find an H$\alpha$ equivalent width of 0.236$\pm$0.021 Å, placing
it in the inactive region of the H$\alpha$–age relations of Kiman et al.
(2021), making an accurate age determination difficult. Kiman et al. (2021)
point out that while stars are found in the inactive region at all ages, the
large majority of inactive stars tend to be those with older ages. While we
can not resolve the metallicity variation noted above, the inactivity favors
the lower metallicity and older age. No light curve is available from TESS or
in the literature. Running the full Gaia kinematics of LSPM J0738+5254 through
BANYAN $\Sigma$, we find that it receives a 99.9% field probability.

CWISE J073831.31+525453.7 (CW0738+5254) is a common proper motion companion
with LSPM J0738+5254. While the $J$-band peak of CW0738+5254’s spectrum is
most similar to that of the L4 near infrared standard as shown in Figure 5,
its continuum appears too dissimilar from the standard to type it through
comparison of the J$-$band alone. A band-by-band classification (cf. Marocco
et al. 2014; Cruz et al. 2018) shows that CW0738+5254 is indeed best fit by
the L4 near infrared spectral standard. However, its spectrum is much redder
than typical field dwarfs of the same spectral type. Using the indices of
Allers & Liu (2013) we find that CW0738+5254 receives a gravity score of FLD-G
(Table 5, indicating an object with field age surface gravity. No other
features of youth, such as the distinctive triangular $H$-band, are visible in
CW0738+5254’s spectrum in Figure 5. Looking at its position on the WISE CMD
(Figure 1), using the parallax of its suspected host, we find that CW0738+5254
appears near normal for its spectral type, although slightly fainter. Using an
estimated parallax for CW0738+5254 based on spectral type - absolute magnitude
estimates, we ran the kinematics of CW0738+5254 through BANYAN $\Sigma$,
finding a 5.7% membership probability for the AB Doradus moving group and a
94.3% field probability.

LSPM J0738+5254 and CW0738+5254 are separated on the sky by an angular
separation of $\sim$11″. At the Gaia distance of LSPM J0738+5254, this
translates to a projected physical separation of $\sim$530 AU. As we do not
find any distinguishing characteristics of youth for either CW0738+5254 or
LSPM J0738+5254, we believe LSPM J0738+5254 + CW0738+5254 to be a field age
widely separated M3+L4 (red) binary.

###  C.2 2MASS J09435055+3356550 + CWISE J094352.22+335639.1

2MASS J09435055+3356550 (2M0943+3356) is a high proper motion star with an
estimated phototype of M3.5 (Cook et al., 2016). Looking at 2M0943+3356 on the
Gaia CMD in Figure 1, we see that it appears near normal for its spectral
type. Sprague et al. (2022) finds a near solar metallicity of
[Fe/H]=0.067$\pm$0.007. No other information, such as rotational period,
H$\alpha$, UV, or X-Ray emission is available in the literature. Running the
full Gaia kinematics of 2M0943+3356 through BANYAN $\Sigma$, we find a 99.9%
field probability.

CWISE J094352.22+335639.1 (CW0943+3356) is a common proper motion companion
with 2M0943+3356. The near infrared spectrum of CW0943+3356 in Figure 5 is
matched closest with the L2 near infrared spectral standard. However, the
spectrum of CW0943+3356 in Figure 5 appears to have more flux at longer
wavelengths, giving it a redder color than the spectral standard. Using the
indices of Allers & Liu (2013) we find that CW0943+3356 receives a gravity
score of FLD-G (Table 5). Looking at CW0943+3356 on the WISE CMD in Figure 1,
using the parallax of its suspected host, we see that it appears fainter than
expected for its spectral type, instead occupying a region typically populated
by mid-L type dwarfs. Running the proper motions and estimated distance
through BANYAN $\Sigma$, we find a 0.7% membership probability for the Argus
moving group and a 99.3% field probability.

2M0943+3356 and CW0943+3356 are separated on the sky by an angular separation
of $\sim$28″. At the Gaia distance of 2M0943+3356, this translates to a
projected physical separation $\sim$1,820 AU. With kinematics for both
objects, as well as the gravity sensitive spectral indices of CW0943+3356,
consistent with the field population, we do not believe the reddened spectrum
of CW0943+3356 is due to youth. We therefore classify 2M0943+3356 +
CW0943+3356 as a widely separated (M3.5)+L2 (red) system.

###  C.3 2MASS J13032992+5127582 + CWISE J130329.90+512754.0

2MASS J13032992+5127582 (2M1303+5127) is phototyped as spectral type M8
(Reylé, 2018b). Looking at 2M1303+5127 on the Gaia CMD in Figure 1, we see
that it appears near normal for its spectral type. No additional information
is available in the literature. Running the Gaia proper motions and parallax
for 2M1303+5127 through BANYAN $\Sigma$, we find a 20.4% membership
probability for the Carina-Near moving group and a 79.6% field probability.

CWISE J130329.90+512754.0 (CW1303+5127) is a common proper motion companion
with 2M1303+5127. While CW1303+5127 does have a Gaia DR3 source ID, there is
no available astrometry for this source. The near infrared spectrum of
CW1303+5127 in Figure 5 is best fit by the L2 near infrared spectral standard.
CW1303+5127’s spectrum in Figure 1 appears redder than the spectral standard
however, as can be seen in Figure 5. Using the indices of Allers & Liu (2013)
we find that CW1303+5127 receives a gravity score of FLD-G (Table 5. Looking
at CW1303+5127 on the WISE CMD (Figure 1) using the parallax of its suspected
host, we see that it appears near normal for its spectral type. Running the
proper motions and estimated distance of CW1303+5127 through BANYAN $\Sigma$,
we find a 4.1% probability of membership to the Carina-Near moving group, a
59.9% probability of membership to the Argus moving group, and a 36.0% field
probability.

2M1303+5127 and CW1303+5127 are separated on the sky by an angular separation
of $\sim$5″. At the Gaia distance of 2M1303+5127, this translates to a
projected physical separation of $\sim$240 AU. While both CW1303+5127 and
2M1303+5127 appear to receive an elevated probabilities of belonging to the
Argus or Carina-Near moving groups, the gravity score of FLD-G for CW1303+5127
suggests the possibility of the pair being field age contaminants. Higher
accuracy proper motions for CW1303+5127, as well as radial velocities for both
components, are required in order to confirm any potential memberships to
young moving groups. As the age of this system is ambiguous at this time, we
believe 2M1303+5127 + CW1303+5127 to be a widely separated field age (M8)+L2
(red).

Table 1: High Probability Co-moving Systems.

Name | SpTypea | Sp$\\_$ref | $\mu_{\alpha}$ | $\mu_{\delta}$ | RV | vsin(i) | Plx | Ast$\\_$ref | Sep | Sep | Probability
---|---|---|---|---|---|---|---|---|---|---|---
|  |  | (mas/yr) | (mas/yr) | km s-1 | km s-1 | (mas) |  | $\arcsec$ | (au) | $\%$
TYC 8025-428-1 (0000-4812) | G5 | 8 | 166.67 $\pm$ 0.01 | -54.58 $\pm$ 0.01 | 2.84 $\pm$ 0.16 | 6.85 $\pm$ 3.05 | 11.42 $\pm$ 0.01 | 2 | — | — | 100
CWISE J000021.45-481314.9 | (L4) | 1 | 172.91 $\pm$ 20.50 | -67.03 $\pm$ 19.20 | — | — | — | 3 | 57 | 5,000 |
WDJ002027.83-153546.50 | (WD) | 9 | 173.66 $\pm$ 0.30 | 6.55 $\pm$ 6.64 | $\pm$ | — | 10.52 $\pm$ 0.32 | 2 | — | — | 100
CWISE J002029.72-153527.5 | (M9) | 1 | 173.03 $\pm$ 13.10 | 16.64 $\pm$ 12.50 | — | — | — | 3 | 31 | 3,000 |
UCAC3 113-933 (0021-3346) | (M0) | 1 | 113.50 $\pm$ 0.02 | -51.36 $\pm$ 0.03 | 12.47 $\pm$ 0.52 | 11.22 $\pm$ 5.42 | 19.02 $\pm$ 0.02 | 2 | — | — | 97.7
CWISE J002101.45-334631.5 | (T1) | 1 | 166.95 $\pm$ 60.20 | 84.15 $\pm$ 56.60 | — | — | — | 3 | 116 | 6,100 |
2MASS J00215427-4514590 | (M5) | 1 | 201.94 $\pm$ 0.02 | -99.17 $\pm$ 0.03 | 11.06 $\pm$ 4.26 | — | 18.46 $\pm$ 0.04 | 2 | — | — | 99.7
CWISE J002159.02-451434.4 | (T3) | 1 | 178.56 $\pm$ 68.50 | -75.73 $\pm$ 68.80 | — | — | — | 3 | 54 | 2,900 |
WT 5A (0024-4031) | (M4.5) | 1 | 189.26 $\pm$ 0.07 | 5.84 $\pm$ 0.07 | $\pm$ | — | 18.66 $\pm$ 0.07 | 2 | — | — | 99.9
WT 5B (0024-4031) | (M4.5) | 1 | 185.38 $\pm$ 0.03 | 7.50 $\pm$ 0.03 | -2.85 $\pm$ 4.55 | — | 18.56 $\pm$ 0.03 | 2 | 2 | 10 |
CWISE J002414.30-403053.7 | (L5) | 1 | 159.72 $\pm$ 24.50 | 26.96 $\pm$ 24.34 | — | — | — | 4 | 78 | 4,200 |
CD-26 134 (0026-2608) | K2 | 1 | 277.04 $\pm$ 0.02 | -39.36 $\pm$ 0.01 | 11.25 $\pm$ 0.26 | 7.12 $\pm$ 3.28 | 11.41 $\pm$ 0.01 | 2 | — | — | 100
CWISE J002658.73-260859.6 | (L5) | 1 | 266.72 $\pm$ 23.60 | -12.31 $\pm$ 22.90 | — | — | — | 3 | 29 | 2,600 |
TYC 1744-123-1 (0042+2918) | K7 | 1 | 103.18 $\pm$ 0.02 | -13.28 $\pm$ 0.01 | -8.76 $\pm$ 0.37 | — | 14.61 $\pm$ 0.02 | 2 | — | — | 99.5
CWISE J004218.67+291730.6 | (L5) | 1 | 87.51 $\pm$ 16.40 | -10.54 $\pm$ 15.10 | — | — | — | 3 | 103 | 7,100 |
LAMOST J004950.95+423653.6 | M4 | 1 | -86.20 $\pm$ 0.01 | -45.12 $\pm$ 0.01 | -65.07 $\pm$ 0.95 | — | 17.77 $\pm$ 0.02 | 2 | — | — | 100
CWISE J004945.38+423619.3 | (L6) | 1 | -72.78 $\pm$ 13.70 | -45.96 $\pm$ 13.00 | — | — | — | 3 | 70 | 3,900 |
CD-24 407 (0056-2402) | K2 | 1 | -108.42 $\pm$ 0.02 | -119.61 $\pm$ 0.02 | 18.11 $\pm$ 0.17 | — | 14.79 $\pm$ 0.02 | 2 | — | — | 100
CWISE J005635.48-240401.9 | (L8) | 1 | -139.18 $\pm$ 30.60 | -93.17 $\pm$ 29.20 | — | — | — | 3 | 102 | 6,900 |
LP 466-214 (0104+1337) | M4.5 | 13 | 217.26 $\pm$ 0.08 | -91.50 $\pm$ 0.06 | $\pm$ | — | 15.63 $\pm$ 0.07 | 2 | — | — | 100
CWISE J010424.50+133949.8 | (L8) | 1 | 241.71 $\pm$ 24.70 | -112.57 $\pm$ 25.30 | — | — | — | 3 | 164 | 10,000 |
WDJ011001.84-592642.22 | (WD) | 9 | 212.99 $\pm$ 0.97 | 167.56 $\pm$ 0.94 | $\pm$ | — | 9.11 $\pm$ 0.79 | 2 | — | — | 100
CWISE J011003.61-592650.8 | (L2) | 1 | 189.95 $\pm$ 16.90 | 182.54 $\pm$ 14.90 | — | — | — | 3 | 15 | 1,700 |
LP 883-372 (0127-3003) | M3 | 1 | 129.03 $\pm$ 0.02 | -5.69 $\pm$ 0.02 | 10.83 $\pm$ 0.58 | — | 20.18 $\pm$ 0.02 | 2 | — | — | 99.9
CWISE J012715.52-300246.9 | (T0) | 1 | 103.80 $\pm$ 16.70 | 3.84 $\pm$ 16.60 | — | — | — | 3 | 51 | 2,500 |
CWISE J013716.34+342352.7 | M3 | 1 | 122.33 $\pm$ 0.02 | 0.96 $\pm$ 0.02 | -2.42 $\pm$ 0.81 | — | 16.22 $\pm$ 0.02 | 2 | — | — | 99.1
CWISE J013719.04+342837.6 | (L6) | 1 | 171.69 $\pm$ 18.90 | 11.18 $\pm$ 17.70 | — | — | — | 3 | 287 | 18,000 |
CWISE J015905.30+105543.2 | M3 | 13 | -109.45 $\pm$ 0.07 | -81.41 $\pm$ 0.05 | $\pm$ | — | 9.01 $\pm$ 0.05 | 2 | — | — | 100
CWISE J015905.58+105551.5 | (M9) | 1 | -110.13 $\pm$ 7.56 | -66.70 $\pm$ 8.41 | — | — | — | 5 | 9 | 1,000 |
NLTT 6937 (0205+5944) | (M4) | 1 | 240.61 $\pm$ 0.02 | 64.12 $\pm$ 0.03 | 46.53 $\pm$ 1.58 | — | 21.99 $\pm$ 0.03 | 2 | — | — | 100
CWISE J020538.20+594452.2 | M8$\beta$ | 1 | 241.46 $\pm$ 0.25 | 64.73 $\pm$ 0.36 | — | — | 21.58 $\pm$ 0.32 | 2 | 4 | 190 |
CWISE J022454.10+152633.8 | (T1) | 1 | 75.53 $\pm$ 20.30 | -60.55 $\pm$ 20.50 | — | — | — | 3 | — | — | 99.9
CWISE J022454.80+152629.5 | L6$\beta$ | 1 | 78.92 $\pm$ 10.70 | -66.07 $\pm$ 9.90 | — | — | — | 3 | 11 | 430 |
G 73-59 (0227+0829) | M3 | 1 | 343.22 $\pm$ 0.36 | -118.33 $\pm$ 0.32 | 75.20 $\pm$ 3.76 | — | 29.56 $\pm$ 0.32 | 2 | — | — | 100
CWISE J022737.75+083008.8 | L5 (blue) | 1 | 366.66 $\pm$ 14.36 | -141.31 $\pm$ 13.97 | — | — | — | 4 | 16 | 550 |
CWISE J025638.42-335454.9 | M2 | 1 | -98.45 $\pm$ 0.01 | -100.76 $\pm$ 0.02 | -2.94 $\pm$ 19.86 | 22.83 $\pm$ 19.48 | 17.30 $\pm$ 0.02 | 2 | — | — | 100
CWISE J025645.16-335008.9 | (L2) | 1 | -85.40 $\pm$ 8.80 | -99.81 $\pm$ 7.90 | — | — | — | 3 | 298 | 17,000 |
BPS CS 22963-0014 (0300-0622) | K7 | 11 | 108.24 $\pm$ 0.02 | -62.33 $\pm$ 0.02 | 62.07 $\pm$ 0.49 | — | 14.90 $\pm$ 0.02 | 2 | — | — | 99.9
CWISE J030005.73-062218.6 | (L9) | 1 | 113.35 $\pm$ 24.90 | -97.03 $\pm$ 26.70 | — | — | — | 3 | 63 | 4,200 |
CWISE J032852.72-112345.6 | (M4) | 1 | 142.54 $\pm$ 0.05 | -106.86 $\pm$ 0.05 | $\pm$ | — | 9.36 $\pm$ 0.04 | 2 | — | — | 100
CWISE J032853.32-112332.6 | (L2) | 1 | 123.87 $\pm$ 23.60 | -124.84 $\pm$ 24.40 | — | — | — | 3 | 16 | 1,700 |
UCAC4 204-003910 (0403-4916) | M2 | 1 | 118.69 $\pm$ 0.02 | 161.28 $\pm$ 0.02 | 31.67 $\pm$ 0.97 | — | 23.90 $\pm$ 0.01 | 2 | — | — | 99.8
CWISE J040351.12-491605.4 | (T7) | 1 | 234.49 $\pm$ 65.40 | 167.85 $\pm$ 66.40 | — | — | — | 3 | 33 | 1,400 |
LSPM J0407+1911A | (M4) | 13 | 101.70 $\pm$ 0.07 | -131.37 $\pm$ 0.05 | 30.14 $\pm$ 3.76 | — | 18.77 $\pm$ 0.06 | 2 | — | — | 99.7
LSPM J0407+1911B | (M4) | 1 | 99.05 $\pm$ 0.03 | -137.12 $\pm$ 0.03 | 32.13 $\pm$ 3.54 | — | 18.90 $\pm$ 0.03 | 2 | 2 | 100 |
CWISE J040702.50+190944.6 | (T1) | 1 | 173.55 $\pm$ 57.20 | -167.10 $\pm$ 62.60 | — | — | — | 3 | 298 | 16,000 |
G 192-8 (0555+5105) | M2 | 1 | 364.16 $\pm$ 0.02 | -489.36 $\pm$ 0.01 | 13.85 $\pm$ 1.25 | — | 20.21 $\pm$ 0.02 | 2 | — | — | 100
CWISE J055515.83+510514.0 | M7 | 1 | 362.81 $\pm$ 0.19 | -489.88 $\pm$ 0.14 | — | — | 20.51 $\pm$ 0.18 | 2 | 23 | 1,100 |
HD 40781 (0559-3842) | G0 | 15 | -19.90 $\pm$ 0.44 | 67.94 $\pm$ 0.45 | 24.92 $\pm$ 0.75 | 6.72 $\pm$ 1.53 | 16.49 $\pm$ 0.36 | 2 | — | — | 100
CWISE J055909.00-384219.8 | (L4) | 1 | -29.94 $\pm$ 10.30 | 77.61 $\pm$ 10.60 | — | — | — | 3 | 54 | 3,300 |
WDJ060159.98-462534.40 | (WD) | 9 | 93.99 $\pm$ 0.15 | -140.10 $\pm$ 0.17 | $\pm$ | — | 17.83 $\pm$ 0.14 | 2 | — | — | 100
CWISE J060202.17-462447.8 | (L8) | 1 | 107.78 $\pm$ 13.00 | -158.00 $\pm$ 14.80 | — | — | — | 3 | 52 | 2,900 |
CWISE J061959.56+344631.3 | M4 | 1 | 91.63 $\pm$ 0.03 | -119.65 $\pm$ 0.02 | 43.47 $\pm$ 2.93 | — | 27.28 $\pm$ 0.03 | 2 | — | — | 100
CWISE J062000.01+344641.3 | L4 | 1 | 87.45 $\pm$ 8.31 | -124.88 $\pm$ 7.77 | — | — | — | 4 | 12 | 420 |