A complete accounting of nearby objectsfrom the highestmass white dwarf progenitors down to lowmass brown dwarfsis now possible thanks to an almost complete set of trigonometric parallax determinations from Gaia groundbased surveys and Spitzer followup. We create a census of objects within a Suncentered sphere of 20 pc radius and check published literature to decompose each binary or higherorder system into its separate components. The result is a volumelimited census of 3600 individual star formation products useful in measuring the initial mass function across the stellar lt8M SUBSUB and substellar 5M SUBJupSUB regimes. Comparing our resulting initial mass function to previous measurements shows good agreement above 0.8M SUBSUB and a divergence at lower masses. Our 20 pc space densities are best fit with a quadripartite power law xi MdNdMpropto Malpha   with longestablished values of   2.3 at high masses 0.55 lt M lt 8.00M SUBSUB and   1.3 at intermediate masses 0.22 lt M lt 0.55M SUBSUB but at lower masses we find   0.25 for 0.05 lt M lt 0.22M SUBSUB and   0.6 for 0.01 lt M lt 0.05M SUBSUB. This implies that the rate of production as a function of decreasing mass diminishes in the lowmass starhighmass brown dwarf regime before increasing again in the lowmass brown dwarf regime. Correcting for completeness we find a star to brown dwarf number ratio of currently 41 and an average mass per object of 0.41 M SUBSUB.
['Initial mass function', 'Stellar mass functions', 'Brown dwarfs', 'Trigonometric parallax', 'Solar neighborhood', 'Binary stars', '796', '1612', '185', '1713', '1509', '154', 'Astrophysics - Solar and Stellar Astrophysics', 'Astrophysics - Earth and Planetary Astrophysics', 'Astrophysics - Astrophysics of Galaxies']
The Initial Mass Function Based on the Fullsky 20 pc Census of 3600 Stars and Brown Dwarfs
{'The Initial Mass Function Based on the Full-sky 20-pc Census of ∼ 3,600 Stars and Brown Dwarfs': "J. DAVY KIRKPATRICK, 1, 2 FEDERICO MAROCCO, 1, 2 CHRISTOPHER R. GELINO, 3 YADUKRISHNA RAGHU, 4, 2 JACQUELINE K. FAHERTY, 5, 2 DANIELLA C. BARDALEZ GAGLIUFFI, 5, 2 STEVEN D. SCHURR, 1 KEVIN APPS, 6 ADAM C. SCHNEIDER, 7, 2 AARON M. MEISNER, 8, 2 MARC J. KUCHNER, 9, 2 DAN CASELDEN, 5, 2 R. L. SMART, 10 S. L. CASEWELL, 11 ROBERTO RADDI, 12 AURORA KESSELI, 1 NIKOLAJ STEVNBAK ANDERSEN, 13, 2 EDOARDO ANTONINI, 2 PAUL BEAULIEU, 2 THOMAS P. BICKLE, 14, 2 MARTIN BILSING, 2 RAYMOND CHIENG, 2 GUILLAUME COLIN, 2 SAM DEEN, 2 ALEXANDRU DEREVEANCO, 2 KATHARINA DOLL, 2 HUGO A. DURANTINI LUCA, 2 ANYA FRAZER, 2 JEAN MARC GANTIER, 2 LÉOPOLD GRAMAIZE, 2 KRISTIN GRANT, 2 LESLIE K. HAMLET, 2 HIRO HIGASHIMURA ( 東 村 滉 ) , 15 MICHIHARU HYOGO, 16, 2 PETER A. JAŁOWICZOR, 2 ALEXANDER JONKEREN, 2 MARTIN KABATNIK, 2 FRANK KIWY, 2 DAVID W. MARTIN, 2 MARIANNE N. MICHAELS, 2 WILLIAM PENDRILL, 2 CELSO PESSANHA MACHADO, 2 BENJAMIN PUMPHREY, 2 AUSTIN ROTHERMICH, 17, 2 REBEKAH RUSSWURM, 2 ARTTU SAINIO, 2 JOHN SANCHEZ, 2 FYODOR THEO SAPELKIN-TAMBLING, 2 JÖRG SCHÜMANN, 2 KARL SELG-MANN, 2 HARSHDEEP SINGH, 2 ANDRES STENNER, 2 GUOYOU SUN ( 孙 国 佑 ) , 18, 2 CHRISTOPHER TANNER, 2 MELINA THÉVENOT, 2 MAURIZIO VENTURA, 2 NIKITA V. VOLOSHIN, 2 JIM WALLA, 2 ZBIGNIEW W ˛EDRACKI, 2 JOSE I. ADORNO, 19, 5 CHRISTIAN AGANZE, 20 KATELYN N. ALLERS, 21 HUNTER BROOKS, 22, 2 ADAM J. BURGASSER, 23 EMILY CALAMARI, 5, 24 THOMAS CONNOR, 25, 26 EDGARDO COSTA, 27 PETER R. EISENHARDT, 28 JONATHAN GAGNÉ, 29 ROMAN GERASIMOV, 30 EILEEN C. GONZALES, 31, 32, ∗ CHIH-CHUN HSU, 33, 34 ROCIO KIMAN, 35 GUODONG LI, 36, 37 RYAN LOW, 38 ERIC MAMAJEK, 39 BLAKE M. PANTOJA, 40 MARK POPINCHALK, 41, 42, 24 JON M. REES, 43 DANIEL STERN, 28 GENARO SUÁREZ, 5 CHRISTOPHER THEISSEN, 23 CHAO-WEI TSAI, 36 JOHANNA M. VOS, 44, 5 DAVID ZUREK, 5 AND THE BACKYARD WORLDS: PLANET 9 COLLABORATION  \n1 IPAC, Mail Code 100-22, California Institute of Technology, 1200 E. California Blvd., Pasadena, CA 91125, USA  \n2 Backyard Worlds: Planet 9  \n3 NASA Exoplanet Science Institute, Mail Code 100-22, California Institute of Technology, 770 S. Wilson Avenue, Pasadena, CA 91125, USA 4 Washington High School, 38442 Fremont Blvd., Fremont, CA 94536, USA  \n5 Department of Astrophysics, American Museum of Natural History, Central Park West at 79th Street, New York, NY 10024, USA 6 Independent scholar, UK  \n7  \nUnited States Naval Observatory, Flagstaff Station, 10391 West Naval Observatory Road, Flagstaff, AZ 86005, USA  \n8 NSF's National Optical-Infrared Astronomy Research Laboratory, 950 N. Cherry Ave., Tucson, AZ 85719, USA  \nNASA Goddard Space Flight Center, Exoplanets and Stellar Astrophysics Laboratory, Code 667, Greenbelt, MD 20771, USA  \n10  \nIstituto Nazionale di Astrofisica, Osservatorio Astrofisico di Torino, Strada Osservatorio 20, I-10025 Pino Torinese, Italy  \n11  \nSchool of Physics and Astronomy, University of Leicester, University Road, Leicester LE1 7RH, UK  \nUniversitat Politècnica de Catalunya, Departament de Física, c/ Esteve Terrades 5, 08860 Castelldefels, Spain  \n13 Sygehus Lillebalt, Department of Cardiology, Kolding, Denmark  \n14 School of Physical Sciences, The Open University, Milton Keynes, MK7 6AA, UK  \n15 Earl of March Intermediate School, 4 The Pkwy, Kanata, ON K2K 1Y4, Canada  \n16 Meisei University, 2-1-1 Hodokubo, Hino, Tokyo 191-0042, Japan  \nPhysics Department, University of Central Florida, 4000 Central Florida Boulevard, Orlando, FL 32816, USA  \n18 Xingming Observatory, Mt. Nanshan, Urumqi, 830011, Xinjiang, PR China  \n19 Department of Physics, University of Miami, Coral Gables, FL 33124, USA  \n20 Department of Physics, Stanford University, Stanford CA 94305, USA  \n21 Department of Physics and Astronomy, Bucknell University, Lewisburg, PA 17837, USA  \n22 Department of Astronomy and Planetary Science, Northern Arizona University, Flagstaff, AZ 86011, USA  \nDepartment of Astronomy & Astrophysics, University of California San Diego, 9500 Gilman Drive, La Jolla, CA 92093, USA  \n24 The Graduate Center, City University of New York, New York, NY 10016, USA  \n25  \nJet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Drive, Pasadena, CA 91109, USA  \n26  \nCenter for Astrophysics  \n|  \nHarvard & Smithsonian, 60 Garden St., Cambridge, MA 02138, USA  \n27 Universidad de Chile, Casilla 36-D, Santiago, Chile  \nJet Propulsion Laboratory, California Institute of Technology, MS 169-237, 4800 Oak Grove Drive, Pasadena, CA 91109, USA  \n29 Institute for Research on Exoplanets, Université de Montréal, Montréal, Canada  \n30 Department of Physics & Astronomy, University of Notre Dame, Notre Dame, IN 46556, USA  \nDepartment of Physics and Astronomy, San Francisco State University, 1600 Holloway Avenue, San Francisco, CA 94132, USA  \n31  \n32 Department of Astronomy and Carl Sagan Institute, Cornell University, 122 Sciences Drive, Ithaca, NY 14853, USA  \n9  \n23  \n28  \n12  \n17", 'KIRKPATRICK ET AL.': 'Table A1 (continued)  \n- a The bands refer to 1 = ch1 = 3.6 µ mand 2 = ch2 = 4.5 µ m.  \n- c Source identification is uncertain or blended in both bands.\n- d Source is undetected in both bands, as these data are very shallow.\n- e Source is badly blended and not extracted in ch1.', 'ABSTRACT': 'Acomplete accounting of nearby objects - from the highest-mass white dwarf progenitors down to low-mass brown dwarfs - is now possible, thanks to an almost complete set of trigonometric parallax determinations from Gaia, ground-based surveys, and Spitzer follow-up. We create a census of objects within a Sun-centered sphere of 20-pc radius and check published literature to decompose each binary or higher-order system into its separate components. The result is a volume-limited census of ∼ 3,600 individual star formation products useful in measuring the initial mass function across the stellar ( < 8 M ⊙ ) and substellar ( ≳ 5 MJup ) regimes. Comparing our resulting initial mass function to previous measurements shows good agreement above 0.8 M ⊙ and a divergence at lower masses. Our 20-pc space densities are best fit with a quadripartite power law, ξ ( M ) = dN / dM ∝ M -α with long-established values of α = 2 . 3 at high masses (0 . 55 < M < 8 . 00 M ⊙ ) and α = 1 . 3 at intermediate masses (0 . 22 < M < 0 . 55 M ⊙ ), but at lower masses we find α = 0 . 25 for 0 . 05 < M < 0 . 22 M ⊙ and α = 0 . 6 for 0 . 01 < M < 0 . 05 M ⊙ . This implies that the rate of production as a function of decreasing mass diminishes in the low-mass star/high-mass brown dwarf regime before increasing again in the low-mass brown dwarf regime. Correcting for completeness, we find a star to brown dwarf number ratio of, currently, 4:1, and an average mass per object of 0.41 M ⊙ .  \nKeywords: stars: mass function - brown dwarfs - parallaxes - stars: distances - solar neighborhood - binaries: close', '1. INTRODUCTION': 'The concept of the initial mass function is one of the most fundamental paradigms in astronomy. It embodies the observational evidence for how the universe turns gas into stars and provides an empirical framework on which to test and inform the underlying theory. The initial mass function has far-reaching influence, from providing the cornerstone for galaxy formation scenarios across all cosmic epochs to determining which stellar and substellar populations we see in our own solar neighborhood.  \nDebate continues on whether the initial mass function is variable with time or dependent on environment, but its description over most of the range of stellar masses in the Milky Way is well determined. Bastian et al. (2010) conclude that the initial mass function is universal for hydrogen-burning stars, at least within the measurement errors of most cur-  \nrent observations, and Andersen et al. (2008) specifically conclude that there is no strong evidence for environmentspecific effects at masses above ∼ 30 MJup . However, far less is known about the mass function at the low-mass end. Knowledge in this area tells us the creation ratio between stars and brown dwarfs and enlightens us on whether planetary mass objects formed via star formation are common compared to those formed via protoplanetary disks.  \nIn this paper, we use recent advances in our knowledge of the nearby stellar census to explore in unprecedented detail the field initial mass function. Gaia has helped refine the nearby census down to spectral types of mid-/late-L out to 20 pc (Gaia Collaboration et al. 2020). For colder spectral types, the WISE mission, together with follow-up parallaxes measured by Spitzer, has filled out this census down to earlyY dwarfs (Kirkpatrick et al. 2019, 2021a), with the help of many other ground-based endeavors (e.g., Best et al. 2021). Our understanding of the low-mass end is dominated by solivagant L, T, and Y dwarfs, but much less is known about the  \n42  \nfrequency with which these low-mass objects exist as companions to hotter objects in the census. We rectify that gap in our understanding by building a complete census of all objects within 20 pc of the Sun and splitting those systems into their individual components.  \nIn Section 2 we use previous nearby star lists, additions from Gaia, and published or newly discovered objects lacking Gaia astrometry to construct the census of objects in the 20-pc volume. In Section 3 we discuss the format of the compiled census, which includes data on nomenclature, astrometry, spectral types, photometry, radial velocities, multiplicity, masses, and effective temperatures. In Section 4 we discuss the methods used to directly measure masses. In Section 5 we discuss the fact that some objects in our sample have strong evidence for multiplicity but generally lack sufficient evidence to characterize the mass of the subcomponents, which is a source of uncertainty in our final analysis. In Section 6 we discuss mass estimation for white dwarf progenitors, giants/subgiants, brown dwarfs, young stars, low metallicity stars (subdwarfs), and normal main sequence stars and discuss what role objects labeled as exoplanets play in our analysis. In Section 7 we perform analysis of the brown dwarf initial mass function, and then we mate that to the stellar initial mass function. In Section 8 we discuss the resulting initial mass function over the entire mass range by compar-  \ning our fit of the functional form to other estimates in the literature, and in Section 9 we summarize our conclusions. Auxiliary data and analyses are found in the Appendices. In Appendix A we present photometric, spectroscopic, and astrometric follow-up used to further characterize 20-pc census members and candidates, and in Appendix B we present a list of the "proximal" systems for each constellation.', '2.1. Building the list of 20-pc systems': 'Our starter list for compiling the census of 20-pc systems was the Preliminary Version of the Third Catalog of Nearby Stars (CNS3; Gliese & Jahreiß 1991), which represents the sum knowledge, prior to large-area digital surveys, of stars believed to lie within 25 pc of the Sun. We took all objects in CNS3 and cross-identified them with the Gaia Early Data Release 3 (eDR3; Gaia Collaboration et al. 2020) to provide updated parallaxes. Objects with parallax values < 50 mas were removed from further consideration, and those with values ≥ 50 mas or lacking a Gaia eDR3 parallax were retained. Two objects listed in the CNS3 as possibly being within 20 pc had no parallax in Gaia DR2, Gaia eDR3, or the literature. These were added to a list, shown in Table 1, of potential 20pc members to consider further. Other additions to this list are discussed in Section 2.1.1.  \nTable 1 . Stars Lacking Trigonometric Parallaxes but Possibly within 20 pc  \nTable 1 continued', '2.1.1. Other published M dwarfs': 'To better complete the M dwarf list, we first consulted the all-sky compilation of Finch et al. (2014), who used the US Naval Observatory fourth CCD Astrograph Catalog (UCAC4; Zacharias et al. 2013) in concert with the American Association of Variable Star Observers (AAVSO) Photometric All-Sky Survey (APASS 5 ) and Two Micron All-Sky Survey (2MASS; Skrutskie et al. 2006) to identify objects within 25 pc of the Sun. The methodology used a suite of color to absolute magnitude relations to provide distance estimates for detections, although this was supplemented with proper motion detection in order to further distinguish nearby stars from background sources. We took this list (their tables 5 and 6) and selected those candidates having Finch et al. (2014) estimated distances ≤ 20 pc and, if available, other published distance estimates ≤ 20 pc from their table 6. This resulted in 267 objects not already in our master census created above. Of these, 251 had parallaxes in Gaia eDR3 (or Gaia DR2, if parallaxes were lacking in eDR3) placing them outside of 20 pc. Of the remaining 16 objects, seven were found to have other published parallaxes or additional distance estimates placing them beyond 20 pc. The final nine possible additions are listed in Table 1 for further scrutiny.  \nSecond, we cross-checked our master table against a volume-complete subsample of 0.1-0.3 M ⊙ Mdwarfs within 15 pc of the Sun (Winters et al. 2021) whose parallax data were pulled from both Gaia DR2 and the literature. We found that all of the host stars in those systems were already included in our master list.  \nThird, we combed through The Solar Neighborhood series of papers by RECONS - specifically papers I (Henry et al. 1994) through XLIX (Vrijmoet et al. 2022) - to identify all objects verified or suspected to fall within 20 pc of the Sun. Our earlier checks had identified all of the confirmed 20-pc objects, but there were, however, a small number of nearby candidates from Winters et al. (2015) that still lack a trigonometric parallax from any source. These were also added to Table 1.  \nAs discussed in the footnotes to Table 1, we have used available photometry and spectroscopy to update the distance estimates for these objects. After additional scrutiny, we find that only one of these - LP 822-37 AB - likely falls within 20 pc, so it has been added to our master census.', '2.1.2. M, L, T, and Y dwarf discoveries from Backyard Worlds': 'Since the recent publication of our 20-pc L, T, and Y dwarf census (Kirkpatrick et al. 2021a), new nearby low-mass stars and brown dwarfs have continued to be recognized via discovery and/or additional follow-up. Examples are a new parallax confirming the nearby nature of the extreme T subdwarf WISEA J181006.18 -101000.5 (Lodieu et al. 2022), the discovery and confirming parallax of the late-T dwarf VVV J165507.19 -421755.5 (Schapera et al. 2022), and the discovery and established physical companionship of the possible Y dwarf companion to Ross 19 (Schneider et al. 2021). Some other 20-pc suspects, such as CWISE J061741.79+194512.8 AB (Humphreys et al. 2023), have also been shown to fall outside the 20-pc volume after additional follow-up. Still other candidates - other isolated field brown dwarfs identified by the Backyard Worlds team - may yet prove to be new members of the 20-pc census.  \nTo assess the status of each of these, we list in Table 2 all newer M, L, T, and Y dwarf discoveries that had initial distance estimates of < 25 pc. To obtain more informed distance estimates of these candidates, in addition to providing additional data on other objects previously believed to be in the 20-pc census, we have searched photometric archives for additional data longward of 1 µ m (along with Gaia magnitudes in the case of brighter sources) and have performed other photometric 6 , spectroscopic, or astrometric follow-up on selected targets. Our own 1.25 µ m and 1.65 µ m followup and reductions, along with our reduction of archival data at 3.6 µ m 4.5 µ m, are described in section A.1. Our optical and near-infrared spectroscopic follow-up is discussed in Section A.2. Additional parallactic measurements are described in Section A.3.  \nUsing this set of compiled data, we have recomputed distance estimates, as listed in Table 2. Column d J is the distance estimate derived by comparing the measured JMKO magnitude to the predicted MJMKO magnitude derived from the M JMKO vs. JMKO -W2 relation of Kirkpatrick et al. (2021a) 7 . This estimate is valid only for objects with JMKO -W2 ≥ 4 . 0 mag, as smaller values may lead to nonunique solutions for M JMKO (figure 20a of Kirkpatrick et al.  \n2021a). Column d H is the distance estimate derived by comparing the measured H magnitude to the predicted MH magnitude derived from the M H vs. spectral type relation of Kirkpatrick et al. (2021a) 8 . The published relation is restricted to types of L0 and later. Column d ch 2 is the distance estimate derived by comparing the measured ch2 magnitude to the predicted Mch 2 magnitude derived from the M ch 2 vs. ch1 -ch2 relation of Kirkpatrick et al. (2021a). This estimate is valid only for objects with 0 . 3 ≤ ch1 -ch2 ≤ 3 . 7 mag, as shown in Figure 18c of Kirkpatrick et al. (2021a). Method 1 (M = 1 in the table) takes the average of these three independent distance measurements - or as many of these as can be derived - as the adopted distance.  \nColumn d W 2 is the distance estimate derived by comparing the measured W2 magnitude to the predicted MW 2 magnitude derived from the M W 2 (M ch 2) vs. W1 -W2 relation of Kirkpatrick et al. (2021a). This estimate is valid only for objects with 1 . 0 ≤ W1 -W2 ≤ 4 . 5 mag, as shown in figure 19c of Kirkpatrick et al. (2021a). Method 2 (M = 2 in the table) takes this estimate as the adopted distance.  \nColumn d G is the distance estimate derived by comparing the measured Gaia G magnitude to the predicted MG magnitude derived from an M G vs. G -J relation derived specifically for this paper. This estimate is valid only for objects with 1 . 5 ≤ G -J ≤ 5 . 0 mag. Method 3 (M = 3 in the table) uses this as the adopted distance. Method 4 (M = 4 in the table) is exactly the same as Method 3 except that its distance estimate, d GRP uses the Gaia G RP magnitude instead of G and uses an M GRP vs. G RP -J relation also derived specifically for this paper. Method 5 (M = 5 in the table) uses the Gaia DR3 parallax, if available, to establish the distance.  \nWhen none of the five estimation methods above apply, we use combinations of colors to solve for degeneracies among possible spectral type or absolute magnitude solutions, as discussed in the notes to the table. For a very small number of objects, the adopted distance is left blank, as no estimate will be possible until additional follow-up is acquired.  \nFinally, as another arbiter of proximity to the Sun, Table 2 lists the measured CatWISE2020 proper motions (Marocco et al. 2021; Eisenhardt et al. 2020) and how significantly those measurements differ from zero. Also, using the color vs. spectral type relations given in Kirkpatrick et al. (2021a), we have estimated spectral type based on the W1 -W2 color (SpW), ch1 -ch2 color (SpS), and J -W2color (SpJW), where the valid color ranges are 0 . 4 ≤ W1 -W2 ≤ 4 . 0 mag, 0 . 3 ≤ ch1 -ch2 ≤ 3 . 0 mag, and 4 . 0 ≤ J -W2 ≤ 8 . 5 mag.  \nOf the 211 candidate objects in the table, 44 have adopted distance estimates placing them closer than 20 pc (res = in, as listed in the table). Although we have tentatively added these 44 objects to the 20-pc census, obtaining parallaxes of all objects in Table 2 still believed to be within 25 pc would', '20-PC MASS FUNCTION': 'Table A1 . Ancillary Spitzer Data  \nTable A1 continued', '2.2. 20-pc stars with newly discovered companions': "While assembling our nearby census, we discovered a small number of objects that fall in close proximity to other, higher mass stars in the list. These known stars and their possible companions are discussed further below and are illustrated in Figure 1.  \nHD 13579 (0215+6740), a K2 dwarf (Bidelman 1985) at 18.6 pc (Gaia DR3) : The motion object CWISER J021550.96+674017.2 from Table 2 was discovered by D. Caselden during a targeted search for companions to known 20-pc stars using multi-epoch imaging data from WISE (Figure 1a). Follow-up JMKO -band photometry from Keck/MOSFIRE (see Section A.1) shows that current location of the companion coincides with a background source, rendering the J -W2 = 0.78 ± 0.02 mag color useless as a gauge of spectral type. The measurement of W1 -W2 = 0.49 ± 0.02 mag from the CatWISE2020 Reject Catalog (Marocco et al. 2021) is also contaminated, as the WISE imaging sequence shown in WiseView 9 (Caselden et al. 2018) indicates that this is a much redder source. The motion measurement from the CatWISE2020 Reject Catalog, 392 ± 11 mas yr -1 in RA and -102 ± 10 mas yr -1 in Dec, is also contaminated by background sources but shows a magnitude and direction roughly similar to the values for the 41 '' -separated K2 star HD 13579 (518.178 ± 0.012 mas yr -1 in RA and -305.636 ± 0.014 mas yr -1 in Dec; Gaia DR3). Using the WISE W2 epochal positions from the unTimely Catalog (Meisner et al. 2023a), a linear least-squares fit results in motions of 698 ± 35 mas yr -1 in RA and -500 ± 71 mas yr -1 in Dec, which are discrepant from the primary's motion values by 5.1 σ and 2.7 σ in RA and Dec, respectively. Curiously, the CatWISE2020 and unTimely motions bracket the Gaia motion values of the primary despite the fact that both the CatWISE2020 and unTimely measurements are WISE-based and are affected by the same background contaminants. Implanting a fake source into the WiseView image sequence with the same W2 magnitude as CWISER J021550.96+674017.2 but with the Gaia-measured motions of HD 13579 provides an excellent match to observed motion of CWISER J021550.96+674017.2 itself, but suggests that the CatWISE2020 value of W2 = 13.84 ± 0.01 mag may be too bright. Given the close separation between the CWISER source and HD 13579 and motions that appear similar, we tentatively denote these as a physical pair with an apparent separation of 760 AU. If associated, the distance to HD 13579 implies a spectral type of > T4.5 for CWISER J021550.96+674017.2 based on the CatWISE W2 magnitude's possibly being biased too bright.  \nHD 17230 (0246+1146), a K6 dwarf (Gray et al. 2003) at 16.2 pc (Gaia DR3) : K. Apps (see Section 3.6.3) notes that there is a fainter star, Gaia DR3 25488745411919488 (G =  \n15.62 mag, ∆ G = 7.51 mag), 3 . '' 6 south of HD 17230 that has no parallax or proper motion solution in Gaia DR3. A search of the Keck Observatory Archive 10 by C. Gelino reveals two epochs of observations of HD 17230 with Keck/NIRC2 behind the adaptive optics system (Wizinowich et al. 2000). Raw images in the Kp and J filters with HD 17230 under a coronagraph (PI: J. Crepp; Program ID: C182N2) clearly show a star located ∼ 3.7 '' from HD 17230 at a position angle of ∼ 195 · . This observation, taken on 2011 Aug 30 UT, can be compared to another taken on 2014 Oct 13 UT (PI: J. Crepp; Program ID: N100N2) in the narrow-band K -continuum. Given the substantial proper motion of HD 17230 of 263.88 ± 0.03 mas yr -1 in RA and -211 . 58 ± 0 . 03 mas yr -1 in Dec (Table 4), the fainter star should fall at a separation of 3 . '' 4 and position angle of 211 · if it were a background source. However, the second epoch shows the secondary at nearly the same separation and position angle as the first epoch, proving that the two stars are a common motion pair. This conclusion is further bolstered by the 2016-epoch Gaia DR3 positions, that place the fainter star at a separation of 3 . '' 64 and position angle of 195 . 3 · from HD 17230. (Figure 1b shows a first-epoch coronagraphic image in which the A component is seen only via its scattered light.) Using the Gaia DR3 parallax for the primary, this implies MG = 14 . 58 mag (spectral type ∼ M8) for the secondary. However, the true type might be somewhat later than this, as the lack of an astrometric solution in Gaia DR3 may mean that this companion is itself a multiple system. This companion, at an apparent physical separation of 59 AU, may be responsible for the radial velocity acceleration seen for HD 17230 over a decades-long timespan by Rosenthal et al. (2021).  \nG 43-23 (1002+1149), an M4 dwarf (Reid et al. 1995) at 17.9 pc (Gaia DR3) : The motion object WISEU J100241.49+145914.9 from Table 2 was discovered by D. Caselden during a targeted search for companions to known 20-pc stars using multi-epoch imaging data from WISE (Figure 1c). This object lies only 15 . '' 6 away from G 43-23, which has astrometry from Gaia DR3 of ϖ abs = 56 . 01 ± 0 . 11 mas, µ α = 157 . 02 ± 0 . 11 mas yr -1 , and µ δ = -235 . 65 ± 0 . 10 mas yr -1 . WISEU J100241.49+145914.9 itself is not listed in either the CatWISE2020 Catalog or the CatWISE2020 Reject Table, but a linear least-squares fit to its epochal unTimely positions (Meisner et al. 2023a) in W2 gives motions of 164 ± 56 mas yr -1 in RA and -233 ± 43 mas yr -1 in Dec, nearly identical to the Gaia motions for the primary. Implanting a W2 = 14.55 mag source with the motions of G 43-23 into the WISE image sequence of WiseView (Caselden et al. 2018) makes for a convincing doppelgänger to WISEU J100241.49+145914.9 itself. The WISEU source's UHS detection at JMKO = 18 . 18 ± 0 . 05 mag results in a color of J -W2 = 3.63 ± 0.11 mag, suggesting a type of ∼ T8.5 and a distance of ∼ 14.3 pc, which is slightly closer than the 17.9 pc distance measured for G 43-23. Nonetheless, given the proximity of the two objects to each other and their nearly identical mo-  \nFigure 1. Images illustrating the six new multiple systems discussed in the text. Each image has north up and east to the left, and the size of each (in arcsec) is noted in the legend. (a) HD 13579 and CWISER J021550.96+674017.2, (b) HD 17230 and Gaia DR3 25488745411919488, (c) G 43-23 and WISEU J100241.49+145914.9, (d) HD 170573 and CWISE J183207.94 -540943.3, (e) G 155-42 and CWISE J184803.45 -143232.3, and (f) 2MASS J19253089+0938235 A and B. All panels show WISE W1+W2 images from WiseView, except for panel (b), which shows a Keck/NIRC2 Kp -band image, and panel (f), which shows a Keck/NIRC2 K -band image.  \n<!-- image -->  \nis to be a physical pair at an apparent physical separation of 280 AU.  \nHD 170573 (1833 -5415), a K4.5 dwarf (Gray et al. 2006) at 19.1 pc (Gaia DR3) : The T7 dwarf CWISE J183207.94 -540943.3 was discovered by G. Colin and B. Pumphrey and first published in Kirkpatrick et al. (2021a), where Spitzer astrometric monitoring gave ϖ abs = 57 . 0 ± 4 . 3 mas, µ α = -129 . 1 ± 11 . 6 mas yr -1 , and µ δ = -172 . 1 ± 9 . 7 mas yr -1 . In assembling the full 20-pc census for this paper, it was noted that this object lies 10 . ' 3 from the K4.5 dwarf HD 170573 (Figure 1d), which has Gaia DR3 astrometric values of ϖ abs = 52 . 29 ± 0 . 02 mas, µ α = -121 . 05 ± 0 . 02 mas yr -1 , and µ δ = -142 . 04 ± 0 . 02 mas yr -1 . Until more accurate astrometry for the T dwarf becomes available, we will consider this pair to be physically associated because these values are only 1.1 σ , 0.7 σ , and 3.1 σ different for ϖ abs , µ α , and µ δ , respectively. If a true binary, the projected separation is 11,800 AU.  \nG 155-42 (1848 -1434), an M3 dwarf (Gaidos et al. 2014) at 17.1 pc (Gaia DR3) : The motion object CWISE J184803.45 -143232.3 from Table 2 was discovered by S.  \nGoodman while searching for unpublished motion objects in WISE imaging data. While assembling the 20-pc census for this paper, it was noted that this source falls 2 . ' 45 away from G 155-42 (Figure 1e). The CatWISE2020 Catalog (Marocco et al. 2021) lists motions for CWISE J184803.45 -143232.3 of µ α = -145 ± 33 mas yr -1 and µ δ = -104 ± 37 mas yr -1 . A linear least-squares fit to the WISE W2 epochal positions from the unTimely Catalog (Meisner et al. 2022) results in motions of µ α = -181 ± 22 mas yr -1 and µ δ = -158 ± 31 mas yr -1 . The Gaia DR3 astrometry for G 155-42 is ϖ abs = 58 . 60 ± 0 . 02 mas, µ α = -236 . 45 ± 0 . 02 mas yr -1 , and µ δ = -237 . 26 ± 0 . 02 mas yr -1 . The measured motion values between the two sources differ by 2.8 σ and 3.6 σ in RA and Dec, respectively, for the CatWISE2020 motion of the potential secondary and by 2.5 σ and 2.6 σ for the unTimely motion. The J -W2color of CWISE J184803.45 -143232.3 from Table 2 suggests a ∼ T7.5 dwarf at a distance of ∼ 15.7 pc, which is sufficiently close to the 17.1 pc distance of G 15542 that we tentatively consider them to be a physical pair with apparent physical separation of 2500 AU, pending improved astrometry for the secondary.  \n2MASS J19253089+0938235, an M8 dwarf (West et al. 2015) at 17.0 pc (Gaia DR3) : C. Gelino finds two epochs of Keck/NIRC2 data for this object in the Keck Observatory Archive. The first epoch (2019 May 22 UT; PI: Bond; Program ID: H299) shows two objects separated by 194 mas at a position angle of 146 · and magnitude difference of ∆ K =0.29 mag. Two objects are still present in the second epoch (2020 Jun 2 UT; PI: Mawet; Program ID: C249) but with a separation of 199 mas and position angle of 137 · (Figure 1f). We conclude that 2MASS J19253089+0938235 is a closelyseparated binary showing orbital motion because the pair shows measurably different separations and position angles but the astrometry of the second object is inconsistent with the motion of a background star, which would have exhibited a relative motion of approximately -80 mas in RA and +240 mas in Dec. Using a UKIDSS Galactic Plane Survey DR11PLUS star visible in the field and located at J2000 RA = 291.3801844 deg and Dec= +9.6377532 deg, we find K =10.53 ± 0.03 mag for 2MASS J19253089+0938235A (the northwest component) and K =10.82 ± 0.03 mag for 2MASS J19253089+0938235B (the southeast component). This object has been flagged as a possible member of the AB Dor Moving Group (Gagné & Faherty 2018).", '2.3. Checks against the Fifth Catalog of Nearby Stars': 'After we had completed our accounting of the 20-pc census, we were presented with an additional opportunity to further check for omissions or subtractions. Golovin et al. (2022) recently published the Fifth Catalog of Nearby Stars (CNS5), a compilation of all stars and brown dwarfs within 25 pc of the Sun. Within the CNS5, there are 3,002 objects with parallaxes of 50 mas or greater, whereas our list has 3,588 individual objects that meet this criterion 11 . For the purposes of checking the completeness of our own census, we find that only twenty-two of these CNS5 objects were not included in our list. These are given in Table 3. Five of these are Gaia discoveries with relatively large Gaia parallax uncertainties. We show in section A.2 that three of these are background objects based on their spectra, and we assume that the other two, given their even larger parallactic errors, are also background objects. Another fifteen have preferred parallaxes that place them beyond 20 pc, and these preferred parallaxes are either revised values in Gaia DR3 or published parallaxes (or new parallaxes discussed in Section A.3) with  \nsmaller uncertainties than those quoted in CNS5 12 . The remaining two objects in Table 3 deserve special note. The first, 2MASSI J0639559 -741844, has a CNS5 parallax with a 16% uncertainty, so we consider our spectrophotometric distance estimate, which places the object beyond 20 pc, to be preferable. (See Kirkpatrick et al. 2021a for a discussion on the credibility of parallaxes when the uncertainties exceed 12.5%.) The second, APMPM J2330 -4737 B, is a bit of a mystery, as we can find no corroborating evidence in the literature that it exists, and this is why it is not included in our census. In conclusion, our comparison to the CNS5 results in no new additions to our list.', '3. THE 20-PC CENSUS': 'Our final 20-pc census is presented in Table 4. The content of this table is described in more detail in the subsections that follow.', '3.1. Nomenclature': 'Not all researchers refer to the same star by the same name, so having a list of aliases is needed. As we entered each object into the census, we searched SIMBAD for alternative names. The name listed under the heading "DefaultName" in Table 4 is the one that appeared as the default name in SIMBAD 13 when our initial search was performed. For all of these objects, a deep dive into the literature is required to establish the current knowledge of multiplicity, spectral type, etc., so we also list alternative names to aid the literature search. Table 4 therefore lists common names (e.g., Sirius), Bayer and Flamsteed designations, and designations from the HR, HD, BD, CD, and CPD catalogs. Table 4 also lists designations from proper motion catalogs (Wolf, Ross, L, LP, G, LHS, LFT, NLTT, LTT, LSPM, SCR, UPM, APMPM, LEHPM, WT, SIPS, PM, and PM J), white dwarf catalogs (WD, LAWD, EGGR), all-sky photometric catalogs (2MASS, WISE), all-sky astrometric catalogs (Gaia, HIC, HIP, TYC, UCAC4, TIC), along with a few other miscellaneous catalogs that also have high usage (GJ, V*, Karmn, **). The field "VarType" is filled with the type of variability seen, if the object is a known variable star; this informa-', 'Table 4 (continued)': 'Table 4 continued', '3.2. Astrometry': 'For each object in the 20-pc census, we list approximate sexagesimal Right Ascension (RA) and Declination (Dec) coordinates at equinox J2000, given under "SexagesimalRA" and "SexagesimalDec" in Table 4. For close multiple systems, the positions of the two objects may be identical, as these coordinates are meant to provide only a crude position for matching the system across catalogs. For more precise coordinates, we also provide RA and Dec in decimal degrees ("RA" and "Dec") at the yearly epoch provided in the "Epoch" column, along with the coordinate uncertainties ("RA\\_unc" and "Dec\\_unc"). Also listed are the absolute parallax ("Parallax") and its uncertainty ("Parallax\\_unc") and the (usually) absolute proper motions and their uncertainties in RA and Dec ("PMRA", "PMDec", "PMRA\\_unc", and "PMDEC\\_unc"). The reference for these decimal coordinates, parallax, and motion measurements is also given ("PlxPMRef"). Note that for some multiple systems, this more precise astrometry may exist only for the composite system or primary and not for each individual component. (As asterisk in the "PlxPMRef" column indicates that the parallax and motion values for another object in the system are used in lieu of actual measurements for this component.) Furthermore, for some recent brown dwarf discoveries, only positions and proper motions are given, as parallaxes have not yet been measured.  \nAs a final note on positions, we provide the constellation in which each object is located ("Constellation"), based on the VizieR tool 17 that uses the constellation boundaries provided by Roman (1987). This column can be used to determine the nearest object in each constellation, as further explored in Appendix B.', '3.3. Spectral types': 'For higher mass stars - typically those with types earlier than mid-M - our primary sources for spectral types were the NStars papers by Gray et al. (2003) and Gray et al. (2006). This was done to assure that as many of our referenced types as possible were classified against a homogeneous system of standards, in this case, the MKK System of Morgan et al. (1943). This system was subsequently updated to the MK system of Johnson & Morgan (1953), which itself was expanded and updated by Morgan & Keenan (1973) (the revised MK system), Keenan & McNeil (1976), and Morgan et al. (1978). (See Hearnshaw 2014 more a more detailed history.)  \nClassification for objects of later type has followed the precepts of the MK System, thereby pushing this homogeneity into the late-M (Boeshaar 1976, Boeshaar & Tyson 1985, Kirkpatrick et al. 1991, Kirkpatrick et al. 2010), L (Kirkpatrick et al. 1999, Kirkpatrick et al. 2010), T (Burgasser et al. 2006, Kirkpatrick et al. 2010), and Y (Cushing et al. 2011, Kirkpatrick et al. 2012) dwarf sequences. Classification is  \ndependent upon the wavelength range over which the typing is done, so Table 4 specifies whether the spectral type was obtained in the visible to photographic near-infrared region ( < 1 µ m; "SpecTypeOpt") or the classical near-infrared region (12.5 µ m; "SpecTypeNIR"). References for the spectral types can be found under "SpTOpt\\_ref" and "SpTNIR\\_ref". For ease of plotting, the spectral types have been converted into a numerical code, with the luminosity type (if listed) ignored. The scale 18 is set so that 0=A0, 10=F0, 20=G0, 30=K0, 40=M0, 50=L0, 60=T0, and 70=Y0; a type of L8.5 would thus be encoded as 58.5. These codes can be found under "SpTOpt\\_indx" and "SpTNIR\\_indx". Note that the original MK classification system\'s standards jump from K5 to K7 to M0 in the late-K sequence, although K6 standards were eventually added in the late 1980s (Keenan & Yorka 1988; Keenan & McNeil 1989). As a result, there are very few objects with codes of ∼ 36 or ∼ 38-39.  \nFor white dwarfs, Table 4 uses types primarily taken from the compilations of Sion et al. (2014) and McCook & Sion (2016), with post-2016 discoveries taken from more recent literature or from Section A.2. The use of these references assures that all white dwarfs are on the spectroscopic classification system proposed by Liebert & Sion (1994). All white dwarf classifications have been assigned based on optical spectra, and the corresponding optical spectral index, "SpTOpt\\_indx", is coded to be the Liebert & Sion (1994) temperature index + 100. That is, our index is set so that DA2=102, DAZ5.8=105.8, DA9.2=109.2, DZ12.6 =112.6, etc. For any white dwarf lacking a temperature index, our spectral index has been arbitrarily assigned a code of 100, as a temperature index of 0.0 cannot exist (and no white dwarf in Table 4 has a temperature index lower than 2.0).', '3.4. Photometry': 'Table 4 provides photometry in several systems that have hemispheric or all-sky coverage. As discussed in section 5.3, objects in the 20-pc census span a vast dynamic range in absolute luminosity, amounting to over twenty-nine magnitudes (a difference of 5 × 10 11 in brightness) in J -band alone. Thus, special care must be taken when choosing photometry for Table 4.  \nFor the traditional "visible" wavelength regime, Gaia eDR3 magnitudes and uncertainties at G , GBP , and GRP are listed ("G", "G\\_BP", "G\\_RP", "G\\_unc", "G\\_BP\\_unc", and "G\\_RP\\_unc"). The brightest reported G -band magnitude is ∼ 2 mag (Gaia Collaboration et al. 2021), and objects with G < 8 mag, GBP ≲ 4 mag, and GRP ≲ 4 mag have residual saturation effects, as detailed in Gaia Collaboration et al. (2021) and Riello et al. (2021). At the faint end, Gaia is complete to G ≈ 20 mag (depending upon source crowding and galactic latitude; Gaia Collaboration et al. 2021), which means  \nthat the more distant late-L dwarfs in the 20-pc census, along with most of the T dwarfs and all of the Y dwarfs, are too faint for Gaia photometric measurements (figure 26 of Gaia Collaboration et al. 2020; see also figure 2 of Theissen 2018).  \nFor the traditional near-infrared wavelength regime, J , H , and K magnitudes are provided, with the caveat that there are two main filter systems in use: the 2MASS filter system 19 and the MKO filter system (Tokunaga et al. 2002). The H -band filter is almost identical between the two, but the J and K filters are quite different. As a result, we provide five separate entries to cover the possibilities JMKO , J 2 MASS , H , K , and Ks - along with their uncertainties ("JMKO", "J2MASS", "H", "K", "Ks", "JMKOerr", "J2MASSerr", "Herr", "Kerr", and "Kserr"). The references for this photometry are given in the "JHK\\_ref" column. The J 2 MASS entries mostly come from 2MASS, whereas the JMKO entries come mostly from surveys based at the United Kingdom Infrared Telescope (UKIRT; e.g., the UKIRT Hemisphere Survey, UHS - McMahon et al. 2013) and the Visible and Infrared Survey Telescope for Astronomy (VISTA; e.g. the VISTA Hemisphere Survey, VHS - Dye et al. 2018). The KMKO entries come primarily from UKIRT-based surveys, whereas the Ks entries come from both 2MASS and VISTA-based surveys. The H entries are pulled from all three sets of surveys. 2MASS provides the only reliable photometry at the bright end of our sample, albeit with large uncertainties, and extends to a S/N = 10 limit of J = 15 . 8 mag, H = 15 . 1 mag, and Ks = 14 . 3 mag at its faint end 20 . UKIRT and VISTA provide reliable photometry between their bright limit ( J ≈ 12 mag for UHS and ∼ 11.512.5 mag in J , H , and Ks for VHS; Dye et al. 2018, GonzálezFernández et al. 2018) and their detection limit ( J ≈ 19 mag for UHS and J ≈ 20 mag, H ≈ 19 mag and Ks ≈ 18 mag for VHS; Dye et al. 2018, González-Fernández et al. 2018) and provide higher angular resolution than 2MASS. We have therefore favored 2MASS photometry for near-infrared magnitudes brighter than ∼ 12 mag and UKIRT/VISTA for fainter magnitudes. For objects even fainter than the UKIRT/VISTA limits, or for objects in areas not yet covered by the public UKIRT and VISTA releases, we have pulled objects from the literature or from Appendix A.1.  \nAt longer near-infrared wavelengths and extending into the near mid-infrared, we also provide WISE-based W1 (3.4 µ m), W2 (4.6 µ m), W3 (12 µ m), and W4 (22 µ m) magnitudes and their uncertainties ("W1", "W2", "W3", "W4", "W1err", "W2err", "W3err", and "W4err"). The reference for the WISE photometry is given in the "WISEphot\\_ref" column. For W1 and W2, magnitudes brighter than W1 ≈ 8 mag and W2 ≈ 7 mag were pulled from the WISE All-sky Source Catalog and fainter magnitudes were pulled from the AllWISE Source Catalog, in accordance with the suggestion  \nmade in the AllWISE Explanatory Supplement 21 . Photometry at W3 and W4 was pulled from the WISE All-sky Release. The only exceptions to the above are objects that were not detected in either of these releases and are instead found only in the CatWISE2020 Catalog. For these sources, only W1 and W2 photometry is presented, as CatWISE2020 has no W3 or W4 photometry.  \nIn the case of 2MASS and WISE photometry, we further provide columns "2MASS\\_contam?" and "WISE\\_contam?". A "yes" in these columns indicates that the associated photometry is likely compromised by another nearby object or artifact, as judged via our by-eye assessments of the multiepoch WiseView image blinks (Caselden et al. 2018), as the poorer image scales of 2MASS (pixel scale of 1 \'\' ) and WISE (pixel size of 1 . \'\' 375) translate to a higher likelihood of source blending.', '3.5. Radial velocities': 'Gaia DR3 provides all-sky radial velocities for stars with GRVS ≲ 14 mag (Katz et al. 2022) and effective temperatures as high as 14,500 K (Blomme et al. 2022). These radial velocities and their uncertainties are also listed in Table 4 ("GaiaRV" and "GaiaRV\\_unc").', '3.6. Multiplicity': 'Even after all systems within the 20-pc volume have been noted, one difficult step remains: correctly determining, based on current knowledge, the multiplicity of each system so that each individual component can be correctly accounted for in the mass distribution. We took a multi-pronged approach at tackling this problem, as described below.', '3.6.1. The "Stellar Ambassadors" program': 'The first approach was to crowdsource the initial reconnaissance of the literature. With the help of the citizen scientist super users of the Backyard Worlds: Planet 9 project, we set up a program whereby volunteers could sign up to investigate the multiplicity of randomly selected 20-pc systems. To make this more enjoyable, the following mission statement was provided:  \n"Our science-fictional Earth Coalition is currently laying the groundwork to explore all of the \'worlds\' within 20 parsecs (65 light years) of the Sun. Scientists on the Earth Coalition\'s Board of Advisors have a list of host \'suns\' within this volume of space, but the details in that list are a bit spotty. The Coalition is seeking to flesh out these details using our Stellar Ambassadors program.  \n"If you choose to become a Stellar Ambassador, your role will be to represent planet Earth to a small number of stellar systems within 20 parsecs. As we reach out for the first time to each of  \nthese stellar neighbors, you will be Earth\'s representative to them. But you need to be knowledgeable of the systems for which you\'re responsible, and that will involve your gaining knowledge of each system you\'re assigned. (By \'system\', we\'re referring to a host star and any of its companions - other stars, brown dwarfs, or exoplanets - in orbit around it.)"  \nEach volunteer was tasked with determining (a) the number of stars, brown dwarfs, and exoplanets in each system, (b) the spectral types of the (sub)stellar components, and (c) the masses of each component, if the masses have been measured. Each Stellar Ambassador was initially assigned a set of ∼ 12 systems, and additional sets would be assigned if the Ambassador wished to analyze more. Importantly, participants were asked to track the reference material that they used for their data collection, regardless of whether they started with SIMBAD, VizieR, Wikipedia, or some other encyclopedic compendium. In total, twenty-one super users participated in the program, which allowed us to cover 56% of the systems with primaries earlier than L0. (All 20-pc objects with primaries later than this had already been scrutinized in Kirkpatrick et al. 2021a.) These efforts were coordinated in weekly and bi-weekly calls with the volunteers.  \nThe product of this exercise was, as expected, an inhomogeneous set of results, as individual Ambassadors concentrated on different portions of the exercise or used entirely different methodologies in their workflows. Nonetheless, it was these varied approaches that enabled us to determine the references on which it would be the most lucrative to focus our early attention. For example, despite the varied approaches, many of the same references kept appearing again and again in the Ambassadors\' reports. These repeating references underscored the vast groundwork laid by exoplanetfinding searches in characterizing potential host stars, as well as the breadth of methods used to measure masses of stars within the solar neighborhood, a topic explored more fully in Section 4. The references that arose from the Stellar Ambassador program were the first resources we used to populate Table 4 with information on multiplicity, mass measurements, and mass estimates.', '3.6.2. In-depth literature checks': 'After this first reconnaissance of the oft-referenced literature, our second approach was the inevitable deep-dive into the literature for each individual object. For this, we used the extensive per-object references compiled by SIMBAD. We concentrated on literature with high-resolution imaging and radial velocity monitoring, in order to judge the multiplicity of each system. We also looked for paper titles that referenced mass measurements and variability (such as eclipsing binaries, RS CVn variables, etc.). Because of time constraints, we were not able to review each reference in detail, but a paper well stocked with results would often allow us to populate Table 4 with information for many systems at once, which sped up the process for objects further down the list.  \nWe also relied heavily on the Washington Double Star Catalog 22 and the Ninth Catalog of Spectroscopic Binaries (Pourbaix et al. 2009), although the former reference lists both confirmed and possible companions that themselves must be studied individually to gauge true companionship.', '3.6.3. The Apps catalog': 'After our in-depth literature checks were completed, we became aware of an unpublished list of objects within 30 pc of the Sun that (now co-author) K. Apps has been carefully curating since 2009. A comparison of the Apps catalog to our list revealed twenty-eight objects, mainly companions, that have been disproved via published literature but that our list still included. These have now been removed from Table 4. The comparison to the Apps list also revealed another twenty-three objects, almost all of which are the second components in spectroscopic binaries or companions revealed by high-resolution imaging, whose discovery literature we had missed. These objects have now been added to Table 4.', '3.6.4. Multiplicity parameters and exoplanets': 'To encapsulate knowledge from the multiplicity checks above, we include several additional columns in Table 4 and split the components of each system into separate rows. An example for one system is illustrated in the mobile diagram (see Evans 1968) of Figure 2. In the column "DefaultName", the entry for the system as a whole appears as " xi UMa AB & WISE 111838.70+312537.9 ". The names of the first subdivision in the mobile diagram of this multiple are denoted by a double hyphen at the beginning of the name, which in this case are " --xi UMa AB " and " --WISE 111838.70+321537.9 ". Further hierarchical branches are denoted by four hyphens (e.g., " ----xi UMa A "), six hyphens (e.g., " ------xi UMa Aa "), etc. The column "#CompsOnThisRow" refers to the number of known components on that row of the table. To select only individual objects in the census, for example, one can downselect only those rows for which "#CompsOnThisRow" equals 1. There is also a "SystemHierarchy" column, giving a code for each division within the system. This is comprised of a four-digit integer (e.g., " 1297 ") that uniquely identifies the system, followed by decimal subdivisions (e.g., " 1297.1 " and " 1297.2 ") to identify subcomponents. For subcomponents that are themselves binaries, further decimal subdivisions (e.g., " 1297.1.1 " and " 1297.1.2 ") are assigned, etc. Table 4 also lists a column called "#CompsInThisSystem" that gives the total number of components in the system. This field is populated only for the top level of each system (those rows having no decimal subdivisions in the "SystemHierarchy" column) and can be summed to find the total number of individual components in the table. Additionally, Table 4 includes a column called "SystemCode" that collapses the "SystemHierarchy" format into an eight-digit integer comprised of the four-digit system identifier followed by  \n<!-- image -->  \nFigure 2. Mobile diagram for the ξ UMa system along with sample columns from Table 4. The mobile diagram at top shows a stylized representation of this quintuple system, illustrating the pair of close doubles ( ξ UMa A and ξ UMa B) and their distant common proper motion companion (WISE J111838.70+312537.9). The table at bottom shows the nine rows for this system, representing the nine vertices (with labels) in the mobile diagram. Table 4 entries for DefaultName, #CompsOnThisRow, SystemHierarchy, #CompsInThisSystem, and SystemCode are shown for illustration.  \nfour additional digits representing any other subdivisions of the "SystemHierarchy" code, but with the decimals removed (e.g., " 12971210 "). Note that when lower subdivisions are lacking, those digits are filled with zeroes. This "SystemCode" column is useful if the user prefers to sort the systems in Table 4 into their mobile diagrams rather than keeping the table\'s default ordering, which sorts by RA.  \nNote that our accounting of components above includes only those stellar and brown dwarf members of the system, but not any of the known exoplanets. For the latter, we also include a column in Table 4 named "#Planets" that reports the number of exoplanets listed in the NASA Exoplanet  \nArchive 23 as of 01 Sep 2022. To match objects from Table 4 to objects in this archive, we used the Transiting Exoplanet Survey Satellite (TESS; Ricker et al. 2015) Input Catalog (TIC; Stassun et al. 2019) designations. It should be noted that, whereas we use a formation-based definition for brown dwarfs in this paper, the NASA Exoplanet Archive uses a mass-based definition for exoplanet vs. brown dwarf and sets the dividing line, somewhat arbitrarily, at 30 M Jup 24 . As a result, there will be some double counting of objects, as these  \nmay appear in both the substellar and exoplanet lists. We will return to this point in Section 6.2.3.', '3.7. Mass parameters and effective temperature': 'The final parameters in Table 4 relate to our need to assign masses to all individual objects within the 20-pc census. In Section 4, we discuss the various methods for which masses can be directly measured. For objects whose masses must, instead, be estimated, Section 6 provides additional discussion. Stars with measured accelerations (see Section 4.1.4) are further discussed in Section 5.1, and objects whose Gaia astrometry suggests hidden companions are discussed in Section 5.2.  \nMass estimation techniques work well for hydrogenburning stars because there is a direct mapping from color, temperature, and spectral type to mass on the main sequence. These same techniques fail for brown dwarfs because color, temperature, and spectral type vary with age, and the age of a brown dwarf is generally an unmeasurable quantity. Estimating the masses for brown dwarfs, therefore, requires a different tack, one that we approach statistically through their distribution of effective temperatures, as further discussed in Section 6.1.3.', '4. MASSES FROM DIRECT MEASUREMENT': 'There are many ways of measuring stellar masses. Some methods (1) measure mass directly using only observational data, (2) lean lightly on theoretical assumptions when a full suite of needed observational data is not available, (3) derive masses by comparing available data to an empirical data grid for stars with directly measured masses, and (4) compare observables to theoretical models. Examples of these third and fourth groups are methodologies such as The Cannon (Ness et al. 2016) and StarHorse (Queiroz et al. 2018). However, the aim of this section is to establish nearby fiducial objects for which masses have been (semi-)directly measured, in order to establish our own empirical grid (method 3) to estimate masses for the remainder of the 20-pc census. Toward this goal, we use the next two subsections to discuss methods 1 and 2 as they have been applied to nearby objects. Table 4 includes directly measured masses for objects that have such values ("Mass" and "Mass err") along with the technique used for the measurement ("Mass method") and its citation ("Mass reference").', '4.1. Multiple systems': 'Mass measurements can be made for objects in binary or multiple systems, once sufficient information has been collected to define the orbits. For compact objects, mass can also be deduced from the gravitational redshift; observationally, this can only be done in multiple systems, as it requires at least one additional, non-compact, co-moving object with which to disentangle the part of the redshift due to radial velocity. More specifics are given below.  \n4.1.1. Visual binaries  \nFor a visual binary whose orbit can be observed, the ratio of the masses is just  \nM 1 M 2 = a 2 a 1 , (1)  \nwhere M 1 and M 2 are the masses of the two objects and a 1 and a 2 are the (physical, not apparent) semi-major axes of their respective orbits. The total mass of the system, M 1 + M 2, can be derived from the equation  \nM 1 + M 2 = 4 π 2 ( a 1 + a 2) 3 GP 2 cos 3 i , (2)  \nwhere G is the gravitational constant, P is the orbital period, and i is the inclination of the orbit on the plane of the sky. The distance to the system must also be measured so that a 1 and a 2 are in physical, not angular, units, and the inclination can be deduced from the difference between the offset of the center of mass and the focus of the projected ellipse (Carroll &Ostlie 1996). Individual masses can be measured by combining equations 1 and 2. A list of visual (and other) multiple systems can be found in the Washington Double Star Catalog 25 .', '4.1.2. Spectroscopic binaries with eclipses': 'For spectroscopic binaries in which the radial velocities of both stars can be measured (SB2s), the ratio of the masses is just  \nM 1 M 2 = v 2 v 1 , (3)  \nwhere v 1 and v 2 are the maximum velocity amplitudes with respect to the mean radial velocity curves of the system. The sum of the masses can be obtained via the equation  \nM 1 + M 2 = P ( v 1 + v 2) 3 2 π G sin 3 i . (4)  \nThe inclination cannot be determined unless the SB2 is also an eclipsing system, in which case the nearly edge-on orientation means that i ≈ 90 · , allowing for a mass determination for both components.  \nThere is a class of eclipsing single-lined spectroscopic binary (SB1) systems for which masses can also be derived (Stassun et al. 2017; Stevens et al. 2018). These are systems with a single stellar host and a transiting exoplanet. Because the combined light of the system is almost entirely that of the host star, available all-sky data sets can provide photometry across a wide swath of the electromagnetic spectrum - from the ultraviolet to the near mid-infrared - so that the star\'s apparent bolometric luminosity can be measured. Accurate parallaxes from Gaia provide the distances needed to convert this to absolute bolometric luminosity. These photometric points span either side of the flux peak in these objects, so they also provide a semi-empirical measurement of effective  \ntemperature, as well. The radius of the host star, R , can then be derived from the Stefan-Boltzmann Law  \nR = √ L bol 4 πσ T eff 4 , (5)  \nwhere L bol is its bolometric luminosity, T eff is its effective temperature, and σ is the Stefan-Boltzmann constant. In the limit where the mass and radius of the exoplanet are far smaller than those of the host star, the density of the host star, ρ , can be calculated directly from observable quantities using the equation  \nρ = 3 π GP 2 a n 3 , (6)  \nwhere a n is the "normalized" semi-major axis (see Sandford & Kipping 2017 for details) and P is the orbital period, both of which can be measured from the transit light curve. (This simplified form assumes a circular orbit. More generalized forms of this equation can be found in Seager & MallénOrnelas 2003.) The stellar mass, M , then follows from  \nM = 4 3 π R 3 ρ. (7)  \nA list of SB1 and SB2 systems (see Pourbaix et al. 2009) can be found at the Centre de Données astronomiques de Strasbourg 26 . A list of 158 detached eclipsing binaries with well measured stellar properties is given in Stassun & Torres (2016).', '4.1.3. Astrometric binaries': 'Astrometric binaries are those systems in which the presence of an unseen companion can be inferred from the nonlinear motion of the primary, once its parallactic motion is accounted for. A careful mapping of the astrometric orbit results in the following measurement  \nM 2 ( M 1 + M 2) 2 / 3 = rap (1 + e ) ( 2 π P √ G ) 2 / 3 (8)  \nwhere M 1 is the mass of the luminous component, M 2 is the mass of the invisible component, rap is the orbital separation of the luminous component at apastron, e is the eccentricity of the orbit, P is the orbital period, and G is the gravitational constant (Andrews et al. 2019).  \nIt is possible to measure individual masses in astrometric binaries if the right conditions are met. We consider here an astrometric binary in which the secondary contributes little or no light to the system, as would be the case in a system comprised of a main sequence star and a black hole, neutron star, cold brown dwarf, or exoplanet companion. In this case, the light of the system comes almost entirely from the primary star, so an analysis of its broad-wavelength spectrum or spectral energy distribution built from broad-wavelength  \nFigure 3. Schematic diagram demonstrating the concept of proper motion anomaly. A binary star system, comprised of an A component (solid black orbit and black points) and a lower-mass B component (dotted pink orbit and pink points) is shown at four separate times corresponding to approximate start and end dates of Hipparcos (left pair) and Gaia DR2 (right pair). The center of mass (grey squares) moves from left to right over time, and the true proper motion of the system over the Hipparcos and Gaia timeframes is represented by the two grey arrows. Assuming that the A component dominates the light of the system, neither Hipparcos nor Gaia will measure this true motion because the center of light will move with component A as the stars orbit their barycenter. The black arrow at left thus shows the proper motion that would be measured by Hipparcos, and the black arrow at right shows the motion measured by Gaia DR2. The disagreement between these two independent measurements is termed "proper motion anomaly" and provides evidence that the system has an unseen component. (For simplicity, we have removed parallactic motion by showing only those points at the same parallax factor, as depicted by the time stamps at the bottom of the figure.)  \n<!-- image -->  \nphotometry can be used to deduce, with the help of empirical relations, its mass, M 1. Then the mass of the companion, M 2, can be measured using Equation 8. Gaia will produce orbits of hundreds of thousands of such astrometric binaries over its anticipated lifetime (Halbwachs et al. 2023).', '4.1.4. Binaries with acceleration (aka proper motion anomaly)': 'Proper motion measurements at two different epochs have the capability of identifying hidden companions if those two motion values differ significantly from one another. (This would be labeled as an astrometric binary, see Section 4.1.3, once additional astrometric epochs are obtained.) The reason is that the proper motion of the system\'s photocenter will deviate from a straight line unless both components contribute equally to the light output. This methodology was first used by Bessel (1844) to deduce hidden companions to Sirius and Procyon. An illustration of the effect, which is known both as "proper motion anomaly" and as "acceleration", is shown in Figure 3. This procedure has seen a recent revival now that high quality Hipparcos motions from the early 1990s and high quality Gaia DR2 motions from the mid-2010s can be compared.  \nThe lack of agreement between the motion measurements is sufficient to identify a hidden companion, and only a few other measurements are needed to derive the companion\'s mass. This can be computed from the following equation  \nfrom Brandt et al. (2019)  \nM = s 2 ( aPM 2 + aRV 2 ) 3 2 G ( ϖ aPM ) 2 , (9)  \nwhere s is the projected separation between the companion and host star, aPM is the host star\'s acceleration on the plane of the sky, aRV is the host star\'s acceleration along the line of sight, and ϖ is the parallax of the system. (See Equation 17 for a different treatment.) This equation holds only if all measurements can be approximated to refer to the same orbital epoch. Otherwise, as detailed in Brandt et al. (2019), more complex orbital fitting is required.', '4.1.5. Compact objects with gravitational redshifts': "Finally, gravitational redshift can be used to measure the mass if the surface gravity can also be determined. Within the 20-pc sample, this is realistically measurable only in relatively massive compact objects like white dwarfs 27 . The observed velocity shift, vgr , due to gravitational redshift is given by  \nvgr = GM Rc , (10)  \nwhere c is the speed of light (e.g., Chandra et al. 2020). Because the star's mass is related to its surface gravity, g , via the equation  \ng = GM R 2 , (11)  \nthe mass can be computed from  \nM = c 2 vgr 2 Gg . (12)  \nThe surface gravity can be measured from the white dwarf's spectrum by comparing to model atmospheres. In practice, though, this method cannot be applied to single white dwarfs because the gravitational redshift is not separable from the radial velocity. If the star is part of a co-moving multiple system or is a member of a young cluster or association, however, then the degeneracy between the radial velocity component and gravitational redshift component can be broken.", '4.2. Single objects': 'Researchers have employed several methods that are capable of measuring the masses of individual objects. These techniques - gravitational lensing, asteroseismology, and surface convection monitoring (aka "flickering") - are described below.', '4.2.1. Gravitational lensing': 'Gravitational lensing occurs when a mass moves very close to the line of sight between an observer and a background object. The mass of the intervening object acts as a lens that alters the apparent position of the background source as seen by the observer (Gaudi 2012) and is potentially measurable for any object. The two temporarily generated images of the background source have a morphology that is azimuthally asymmetric, and this manifests itself observationally as a shift in the centroid. The astrometric shift of the photocenter is given by  \nδ ( t ) = u ( t ) θ E u ( t ) 2 + 2 (13)  \nwhere θ E is the angular Einstein radius, which can be expressed as  \nθ E 2 = 4 GMl c 2 ( Dl -1 -Ds -1 ) (14)  \n(Walker 1995, Lu et al. 2016). Here, u and u represent the scalar and vector time-dependent lens-source separation in the plane of the sky normalized to θ E , Ml is the mass of the lens, and Dl and Ds are the distances to the lens and source, respectively. When the distances to the lens and source are known, the monitoring of the astrometric shift as a function of time enables a measurement of the mass of the lens. These equations show that closer lenses produce larger astrometric signals, which makes this a valuable technique for measuring the masses of nearby objects, the main limitation being that such encounters of a lens and a background source happen only rarely and very accurate astrometry is needed to predict such encounters a priori. This technique has so far been successfully applied to only two objects in the 20-pc sample (Sahu et al. 2017, Zurlo et al. 2018) but promises to become more valuable as more accurate Gaia parallaxes and proper motions become available for stars all across the Milky Way.', '4.2.2. Asteroseismology': "Asteroseismology is the study of oscillations in stellar atmospheres, and the characteristics of these oscillations can be used to deduce a star's physical parameters. Any star having a mechanism that can drive oscillations - such as surface convection, pulsations, tidal effects in a close binary, or opacity effects (the κ -mechanism) - can potentially have its mass measured. Equation 52 in Aerts (2021) can be rewritten to show that the stellar mass, M , can be determined from these oscillations using the relation  \nM ∼ ν max 3 T eff 3 2 ∆ ν 4 . (15)  \n(In the absence of a definitive theoretical model for convection, the scaling of this relation is based on observations of the Sun, as described in Kjeldsen & Bedding 1995.) Here, ∆ ν is the large frequency separation, ν max is the frequency of maximum power, and T eff is the effective temperature. The quantity ν can also be thought of as the inverse of twice the  \nsound travel time between the stellar center and the stellar surface (Eq. 39 of Aerts 2021). Figures 4 and 10 of Aerts (2021) graphically demonstrate how ν and ν max are measured in practice. The effective temperature, T eff, is obtained by comparing broad-wavelength spectroscopy of the star to model atmospheres.", '4.2.3. Surface convection monitoring ("flickering")': "The full asteroseismology treatment above requires highquality data over a sufficient time baseline with which to resolve the individual oscillation modes. However, variations in surface convection alone require less exquisite data and can be used to measure the mass, if certain ancillary quantities have also been well measured (Stassun et al. 2018). The needed quantity is ν max from above, which has been shown to depend on the star's gravity, g , and effective temperature, T eff, through the relation  \ng = ν max √ T eff C (16)  \n(Brown et al. 1991), where C is a normalization constant obtained by calibrating to stars with gravity measurements independently determined via asteroseismology (Kallinger et al. 2016). The effective temperature is, as above, obtained by comparing broad-wavelength spectroscopy to model atmospheres. The mass can then be measured via Equation 11, where the star's radius can be measured directly via interferometry or through the Stefan-Boltzmann Law in Equation 5. The bolometric luminosity can be measured from the aforementioned broad-wavelength spectrum along with an accurate trigonometric parallax.", '5. MULTIPLES LACKING SUFFICIENT DATA FOR MASS DETERMINATION': 'There are some systems for which acceleration has been measured or whose astrometry indicates the presence of multiple components but for which insufficient data exist to compute the masses of the individual objects. Such systems are important to note because their mass accounting is still incomplete. This serves as an additional source of uncertainty in our mass function analysis.', '5.1. Multiples known only through limited acceleration data': 'Currently, there are many accelerating objects within the 20-pc census that lack the additional data needed for companion mass computations via Equation 9. We nonetheless still note these as binaries in Table 4, and we split out those cases here for individual discussion.  \nKhovritchev & Kulikova (2015) have identified likely accelerators by comparing the proper motion measured between the first and second Palomar Observatory Sky Surveys (POSS-I and POSS-II; Minkowski & Abell 1963; Reid et al. 1991; Lasker & STSCI Sky-Survey Team 1998) to a motion derived using first-epoch data from other sky surveys (2MASS, SDSS, WISE) and their own second-epoch follow-up astrometry. With these two independent measurements, they can compare a long-baseline motion over 50 yr to  \nTable 5. New 20-pc Accelerators from the Khovritchev & Kulikova (2015) Sample  \none derived more instantaneously, over only ∼ 10 yr. Brandt (2021) have similarly intercompared the near-instantaneous Hipparcos-measured proper motion from the early 1990\'s, the near-instantaneous Gaia-measured motion from the midto late-2010\'s, and a long-baseline motion constructed from the Hipparcos-to-Gaia baseline. Both sets of authors have identified objects with significant motion discrepancies and labeled these as likely binaries. These objects are noted in Table 4 using the column labeled "Accelerator?".', '5.1.1. Accelerators from POSS vs. recent motion comparison': 'The Khovritchev & Kulikova (2015) list of ∼ 2400 objects covers only a portion of the northern sky (30 · < Dec < 70 · ) for bright ( V < 17 mag), high motion ( µ > 300 mas yr -1 ) stars. Within 20 pc of the Sun, nine such accelerators are identified, only two of which - BD+66 34 and G 96-29 (Capella HL) - were already identified as known multiples in Table 4. The other seven are listed in Table 5.  \nWe note that none of these seven objects is identified as a high-significance accelerator in the Brandt (2021) reference discussed in the following subsection. This is because the Brandt (2021) Hipparcos-to-Gaia accelerations could not be computed for these seven stars, as none are in the Hipparcos Catalog 28 . To further explore the underlying data for these Khovritchev & Kulikova (2015) accelerators, we have produced finder charts that show all seven in the POSS-I, POSSII, 2MASS, and WISE images. A few of these appear to be blended with a background object at one of the POSS epochs. The most notably affected are G 172-30, which is blended at POSS-I with an object fainter by ∆ G = 5.8 mag; Wolf 47 29 , which is blended at POSS-I with an object fainter by ∆ G = 6.4 mag; and G 192-13, which is blended at POSS-II with an object fainter by ∆ G = 6.2 mag. (Ross 10 moves past a star  \nof near-equal magnitude in all of the images, the possible blending being worst at the POSS-II and 2MASS epochs.) This having been noted, whether or not objects with these magnitude differences could perturb the POSS measurements enough to affect the 50-yr proper motion measurements is not clear. Future releases from a longer baseline Gaia data set should determine whether the accelerations seen for these seven objects are real.', '5.1.2. Accelerators from Hipparcos vs. Gaia comparisons': 'The Brandt (2021) list of ∼ 115,000 objects covers the entire sky for objects in common to Hipparcos and Gaia eDR3 ( G ≲ 11 mag). This list also gives the computed χ 2 value between the two proper motions measured with the best precision, which is usually the Gaia-specific and Hipparcos-toGaia measurements. We conservatively set a false alarm rate of Q = e -χ 2 / 2 < 0 . 1%, corresponding to χ 2 > 13 . 8, to select high-confidence accelerators for analysis here. Using this criterion produces ∼ 33,750 objects, of which 194 fall within the 20-pc census. These 194 are denoted in Table 4 with a "yes" in the "Accelerator?" column.  \nKervella et al. (2022) have also produced a catalog of possible accelerators based on a comparison of the short-baseline Gaia-specific motions and the long-baseline Hipparcos-to-Gaia motions. As this list is based on the same underlying data as the list produced by Brandt (2021), many of the same accelerators are flagged by both teams. Under the  \nassumption that the companion mass is much less than that of the primary and that the (circular) orbit is perpendicular to the line of sight, Kervella et al. (2022) have further used the proper motion measures to estimate the mass of the hidden companion using the equation  \nm = ( 4740 . 470 ∆ µ ϖ ) √ rM G , (17)  \nwhere m is the companion mass, M is the primary mass, G is the gravitational constant, r is the orbital radius, ∆ µ (the difference in motion measurements) is in units of mas yr -1 , and ϖ is in units of mas. The constant of 4740.470 is used to convert ∆ µ/ϖ into units of m s -1 (Kervella et al. 2019). Companion masses are estimated using estimated primary masses generally from isochrone fitting for the brightest stars or from an absolute K -band relation for the fainter stars, as described further in Kervella et al. (2022). Companion masses are dependent upon the unknown value of the separation between components, so Kervella et al. (2022) constructed estimates for assumed separations of 3, 5, 10, and 30 AU. In Table 4 we include the extrema of these mass estimates in columns labeled "EstMassAt3AU" and "EstMassAt30AU" for all objects tagged as accelerators. (In a small number of cases, a Brandt 2021 accelerator was not deemed to be an accelerator by Kervella et al. 2022, so these estimates are not given.)  \nTable 6 . 20-pc Accelerators in Known Close Binary/Multiple Systems  \nTable 6 continued  \nTable 6 (continued)  \nTable 6 continued  \nTable 6 (continued)  \nWe divide the resulting list of 20-pc accelerators into three subgroups. The first, listed in Table 6, comprises eightythree objects in known close binary and multiple systems. For all of these, the host star is known to have a close-in companion that Gaia eDR3 fails to detect or provide a full astrometric solution for, and these companions range in mass from the substellar regime into the planetary regime. For a host star at a distance of 10 pc, its Hipparcos-to-Gaia acceleration can be detected if the companion has a separation below a few × 100 AU (Figure 12 from Kervella et al. 2022). Companions at this separation range can also be detected with high-resolution imaging techniques or via radial velocity monitoring, and some have independently measured masses. As one example, the companion in the 19.5-yr spectroscopic binary HD 10307 AB has a measured dynamical mass from Torres (2022) of 0.254 ± 0.019 M ⊙ , and that system has a = 7 . 7 AU, i = 100 · and e = 0 . 44. The Kervella et al. (2022) companion mass estimates of 0.20 M ⊙ at 3 AU and 0.63 M ⊙ at 30 AU bracket the dynamically measured values well, as the assumptions used were reasonable for this system. As another example, the companion to the 1.35-yr spectroscopic binary HD 184467 AB has a measured dynamical mass of 0.868 ± 0.025 M ⊙ (Piccotti et al. 2020), and the system has a = 0 . 7 AU, i = 145 · and e = 0 . 34 (Arenou et al. 2000). The Kervella et al. (2022) companion mass estimate of 0.03 M ⊙ at 3 AU compares unfavorably to the measured  \nvalue possibly because of the Kervella et al. (2022) assumption that the secondary mass is much less than that of the primary. This demonstrates that, although the Kervella et al. (2022) companion mass estimates listed in Table 6 provide a guide as to whether the companion causing the acceleration is already known or is a still hidden member, additional astrometric data is needed before the masses can be reliably measured. As can be seen from the full entries in Table 4, many objects in Table 6 are triples, so it is also unclear how many objects are contributing to the measured acceleration.  \nThe second list, shown in Table 7, gives fourteen objects known to host exoplanets but lacking any "close" stellar or substellar companions. Here we define "close" to mean within ∼ 50 AU. Six of these objects, as listed in the final column of the table, have more widely separated stellar companions at apparent separations of ≳ 70 AU. The Kervella et al. (2022) mass estimates for all fourteen of these objects are quite low and, for assumed separations of a few AU, correspond to masses traditionally thought of as being in the planetary range. Thus, the accelerations for these objects are likely caused by the known exoplanet(s) in the system. Kervella et al. (2022) provides additional analysis on the stars ϵ Eri, Kapteyn\'s Star, ϵ Ind A, and π Men, while noting that Kapteyn\'s Star has no significant proper motion anomaly as measured by them.  \nThe final list, shown in Table 8, has ninety-seven objects whose closest known companions are resolved by Gaia eDR3 or have no known companions at all. For many of these, the nearest known companion falls close enough to the accelerator star ( ≲ 100 AU; Figure 12 of Kervella et al. 2022)  \nTable 7. 20-pc Accelerators whose only Close Companions are Known Exoplanets  \nthat it may be the object causing the acceleration. Examples are CD -44 3045 A, VV Lyn Aa, CD -36 6589 A, Ross 52 A, BD+45 2247 A, and Wolf 1225 A. Objects for which the nearest known companion lies beyond this separation or for which no companions are currently known are the hosts most likely to harbor new additions to the 20-pc census. Examples of stars with likely hidden companions are G 32-7, CD -22 526, HD 13579, LP 837-53, HD 43162 A, HD 52698, G 25034, BD -17 3088, µ Vir, β TrA, and θ Cyg.  \nTables 6-8 highlight that the accounting of all components within the 20-pc census is still incomplete, as there is overwhelming evidence of additional, tightly separated companions. As only < 200 of the ∼ 3,000 Gaia-detected primaries show such evidence, it is tempting to conclude that our tally of higher mass (non-exoplanet) companions is nearing com-  \netion. We caution, however, that our criteria for selecting accelerators was set very conservatively and that many real accelerators likely exist with a measured significance below our cutoff value. As the time baseline of Gaia observations is extended, accelerations will be increasingly sensitive to longer-period companions that, for higher (non-exoplanet) masses, are potentially verifiable with direct imaging techniques. Furthermore, Gaia observations over this same extended time baseline will remove the need to compare to the shallower Hipparcos data, enabling acceleration data for lower-mass primaries between the Hipparcos and Gaia limits (11 ≲ G ≲ 21 mag). Finally, less than a third of all systems in the 20-pc census of Table 4 have both a Hipparcos entry and a Gaia DR3 astrometric solution, so many objects within our sample volume are unavailable for similar acceleration analysis.  \nTable 8 . 20-pc Accelerators with More Distant (or No Known) Companions  \nTable 8 continued  \nTable 8 (continued)  \nTable 8 continued  \nTable 8 (continued)', '5.2. Multiples with large RUWE values': 'The Gaia Renormalized Unit Weight Error (RUWE) is a measure of the goodness of fit of the single-star astrometric model to the observed astrometry and is expected to be ∼ 1.0 if the fit is a good representation (Lindegren et al. 2021b). This parameter is pulled from Gaia DR3 and is listed in the "RUWE" column of Table 4. Values significantly higher than unity can indicate either that the object is an unresolved, physical multiple (Penoyre et al. 2020) or that some other effect is causing the photocenter to deviate from expectations. Two examples of the latter are a chance alignment with a marginally resolved background star or single-star variability that confounds the RUWE renormalization itself (Belokurov et al. 2020). The typically quoted value for selecting likely binaries using this statistic is RUWE > 1.4 (e.g., Fabricius et al. 2021), although Stassun & Torres (2021) have shown that values of 1.0 < RUWE < 1.4 are also highly predictive of unresolved multiplicity. While the RUWE normalization works well across the full population of Gaia-measured stars, Penoyre et al. (2022b) note that it does a somewhat less adequate job when a selection of nearby (d < 100 pc; the GCNS of Gaia Collaboration et al. 2020) stars alone is analyzed. For that reason, they define a new statistic, which they term the Local Unit Weighted Error (LUWE), that improves upon RUWE for these nearer objects.  \nValues of RUWE and LUWE change with each subsequent release of Gaia data, and there is valuable information contained within the differences. Later Gaia releases have data (and astrometric solutions) covering a longer timespan, so for unresolved multiple systems with periods roughly equal to or longer than the timespans of the data release, the RUWE (or LUWE) values may continue to run high or even become larger between Gaia DR2 and Gaia eDR3 simply because the photocentric displacement caused by orbital motion in an unresolved binary makes the single-object astrometric solution fit less well with an extended data set. (See Penoyre et al. 2022a for additional discussion.) Conversely, unresolved binaries with shorter periods should improve and eventually get full astrometric solutions in the Gaia non-single star lists.  \nWith these observations in mind, Penoyre et al. (2022b) devised a two-part criterion to select the most likely hidden multiples in the 100-pc sample: (1) LUWE eDR 3 > 2 and (2) ∆ LUWE ≡ LUWE eDR 3 -LUWE DR 2 > -LUWE eDR 3 / 3. Within our 20-pc census, 104 objects meet these criteria, and these are the ones labeled with a "yes" in the Table 4 column named "LUWE\\_binary?". Of these, 73 are known from previous literature to be binary and were already labeled as such in our census. The other 31, listed in Table 9, are newly identified multiples. Nine of these are part of higher-order multiples, as indicated by the notes in the table. For eight of these systems, Gaia has detections of both the new high-LUWE object and the other component (sometimes a double itself) with which it has physical companionship. The ninth system, however, is a new triple system for which Gaia detects only the new high-LUWE binary G 43-23 but not the common-proper-motion T dwarf compan-  \nWISEU J100241.49+145914.9, discussed in Section 2.2. Note that the LUWE criteria from Penoyre et al. (2022b) are meant to be conservative, so other hidden binaries will exist with LUWE or ∆ LUWE values outside of the bounds noted above.  \nTable 9 . New 20-pc Multiple Systems Identified Through LUWE  \nTable 9 continued  \nTable 9 (continued)', 'NOTE-': "- · (1) 0146 -5339: This is the 35 . ' 9-distant companion to the F9 dwarf q01 Eri.\n- · (2) 0749 -0320: This is the 3 . ' 9-distant companion to the M3.5 binary PM J07498 -0317 AB.\n- · (3) 1002+1459: This object also has a 15 . '' 6-distant companion, WISEU J100241.49+145914.9, announced in this paper (Section 2.2).\n- · (4) 1145 -2021: This is a physical system with the M5e star LP 793-34, 15 . '' 2 distant.\n- · (5) 1417+4525: This is the 59 . '' 2-distant companion of the M0 star BD+46 1951.\n- · (6) 1428+0518: This is a physical system with the M4 star G 65-54, 1 . ' 0 distant.\n- · (7) 1843 -3322: This is a physical system with the M6 star CE 507, 15 . '' 0 distant.\n- · (8) 1936+5013: This is the 1 . ' 9-distant companion to the F3+ dwarf θ Cyg, which is listed amongst the accelerators in Table 8.\n- · (9) 2217 -0848: This is a physical system with the M5 dwarf binary Wolf 1561 BaBb, 7 . '' 9 distant.", '5.3. Multiplicity (and oddities) identified through color-magnitude diagrams': 'In Figures 4-9, we show several color-type, color-color, and color-magnitude diagrams as a final method for identifying unresolved binaries. These diagrams also illustrate the rich diversity of colors and absolute magnitudes that objects within the 20-pc census possess.  \nEach plot shows photometry only for those objects believed to be single components ("# Components this Row" = 1 in Table 4) and whose photometry is uncontaminated ("2MASS\\_contam?" and/or "WISE\\_contam?" not equal to "yes" in Table 4). Each object is color coded by its spectral type, as shown by the color bar in each figure 30 . Preference is given to the near-infrared spectral type if listed; otherwise, the optical spectral type is used. (It should be noted that for stars of type A through M, near-infrared classifications are given in Table 4 only when no optical type is available, so this criterion is only relevant for the L, T, and Y dwarfs.) Each object is plotted as a solid black dot, the center of which is colored if the spectral type is known; that is, objects lacking a spectral type appear only as black dots. Furthermore, for plots that involve J or K bands, preference is given to MKO magnitudes; otherwise 2MASS magnitudes are used.  \nIn Figure 4, only those objects with absolute magnitude uncertainties below 1.0 mag are shown, to keep the plots more legible. In Figure 5, objects are shown only if their uncertainties in MG are below 1.0 mag and their color uncertainties are below 0.5 mag. In Figure 6 (or 7), objects are shown only if their absolute magnitude uncertainties are below 1.0 mag and their J -W2(or H -W2) uncertainties are also below 1.0 mag. In Figure 8, objects are shown only if the color uncertainty is less than 0.5 mag for Gaia-based color plots or less than 1.0 mag for all other colors. In Figure 9, points are shown only if their color uncertainties are generally less than 0.10-0.20 mag.  \nWe have examined each of these diagrams in detail and have identified objects that fall significantly far from the common loci of main sequence stars or white dwarfs to warrant special attention. There are several classes of objects, however, that we do not discuss in this section but address elsewhere: (1) Stars with bright magnitudes may be problematic and have quoted uncertainties insufficiently small to capture these problems. Given that these bright stars are generally well characterized already, we concentrate only on those not believed to be main sequence stars (category #4 below). (2) L, T, and Y dwarfs have already been examined in detail via color-type, color-magnitude, and color-color diagrams in Kirkpatrick et al. (2021a). (3) Subgiant, giant, and bright giant stars are discussed in Section 6.1.2. (4) Low-metallicity (subdwarf) stars are discussed in Section 6.2.2. (5) Young objects are discussed in Section 6.2.1.  \nWe begin with objects whose placement on these diagrams could potentially highlight a problem with their measured parallaxes. These are all objects that the Gaia survey is placing within the 20-pc volume for the first time. With the exception of the object at 11 h 59 m -36 · 34 \' , all of these objects have higher than normal Gaia parallax uncertainties as compared to objects of similar G magnitude. We discuss each of these individually below:  \n- · Gaia EDR3 4966072879648455296 (0229 -3606): This object, whose spectral type has yet to be determined but whose Gaia eDR3/DR3 parallax is 50 . 66 ± 0 . 61 mas, falls near or just above the main sequence on most color-magnitude diagrams. Its apparent magnitudes are similar to 20-pc objects of the same color, so there is no reason to question its inclusion in our census. Its location on color-magnitude diagrams along with its high DR3 RUWE value of 5.438 indicate possible unresolved binarity.\n- · Gaia EDR3 3330473222213987072 (0623+1018): This M3 dwarf (see Section A.2) has a Gaia eDR3/DR3 parallax of 50 . 80 ± 1 . 55 mas. The derived MG value is ∼ 8 magnitudes fainter than that expected for an average M3 dwarf, and the MW 2 value is ∼ 9 magnitudes fainter. The Gaia parallax value for this object is clearly in error, so it has been removed from the 20-pc census.  \nFigure 4. Various absolute magnitudes plotted against spectral type for the 20-pc census. See text for details.  \n<!-- image -->  \nFigure 5. Absolute G -band magnitude plotted against various colors for the 20-pc census. See text for details. The spray of mostly black points (i.e., objects with no measured spectral types) to the left of the main sequence in the GBP -G vs. MG diagram and to the right of the main sequence in the G -GRP vs. MG diagram represents components in close binaries near the Gaia resolution limit. The GBP and GRP magnitudes are calculated from the fluxes in a 3.5 × 2.1 arcsec 2 field, whereas the G magnitudes are calculated from a profile fit to a much higher-resolution image (section 8 of Evans et al. 2018). For binaries just above the Gaia resolution limit, this means that per-component BP and RP fluxes will often include light from the other object, whereas the G flux will not (Halbwachs et al. 2022). This effect pushes such objects blueward in GBP -G color and redward in G -GRP color, as these diagrams show.  \n<!-- image -->  \nFigure 6. Various absolute magnitudes plotted against J -W2 color for the 20-pc census. See text for details.  \n<!-- image -->  \nW2 (mag)  \nFigure 7. Various absolute magnitudes plotted against H -W2 color for the 20-pc census. See text for details.  \n<!-- image -->  \nFigure 8. Various colors plotted against spectral type for the 20-pc census. See text for details.  \n<!-- image -->  \n- · Gaia EDR3 3460907947316392704 (1159 -3634): This is an M9.5 dwarf with a Gaia eDR3/DR3 parallax of 50 . 10 ± 0 . 18 mas. Its apparent magnitudes fall within the range of other M9.5 dwarfs within the 20pc census. In absolute magnitude, it falls above the main sequence by as much as a magnitude for objects of similar color, and its Gaia RUWE value is 1.482. This position on color-magnitude diagrams cannot be explained by binarity alone, but a slightly larger parallax in tandem would solve the discrepancy. In any event, there is no reason to exclude this object from Table 4.\n- · Gaia EDR3 6025146733201615616 (1624 -3212): This object, of unknown type, has a Gaia eDR3/DR3 parallax of 59 . 01 ± 0 . 12 mas. Its apparent magnitude falls in the range expected for objects of similar color within 20 pc. On plots of absolute magnitude vs. color, however, it appears anomalous. On the MG vs. GBP -GRP plot, it falls 0.4 mag more luminous than objects of similar color; on the MG vs. G -J plot, it is also more luminous, but by 2.4 mag. Whether these issues indicate a problem with the measured photometry, the measured astrometry, or both - or whether the object has an unusual spectrum - is currently unknown. This object is retained in Table 4.\n- · Gaia DR2 4062191480304598656 (AB) (1736 -2515): This object, also of unknown type, has a Gaia DR2 parallax of 60 . 24 ± 0 . 83 mas and is a known double (Vrijmoet et al. 2022). Gaia DR3 lists two sources  \nnear this position, but neither have a parallax or proper motion measurement. The apparent magnitude of the DR2 source is at odds with the range expected for objects of similar color within 20 pc for many combinations of apparent magnitude vs. color, such as G vs. G -W2, J vs. J -W2, and H vs. J -Ks . The object also has a very small Gaia-measured proper motion of 25 . 2 ± 1 . 1 mas yr -1 and lies near the Galactic Center at l , b = (2 . · 1 , + 3 . · 1). This object is most likely a background object with a faulty parallax, so we have removed it from the 20-pc census.  \n- · Gaia DR2 1795813379365971072 (2151+2328): No spectrum has been acquired of this object, and it appears to be a very close double in both Gaia DR2 and eDR3/DR3. However, only one of these components has a parallax measurement in DR2, and neither one does in DR3. This object lies well below and blueward of the main sequence on many apparent magnitude vs. color plots such as J vs. G -J , W2 vs. GBP -GRP , and H vs. J -W2. The object also has a small Gaia DR2 motion of only 45 . 7 ± 0 . 7 mas yr -1 . This is likely a background source with a bogus parallax, so we have removed it from the 20-pc census.  \nThe rest of our analysis deals with objects that are outliers for various other reasons. As discussed below, these reasons include possible unresolved binarity, unusual atmospheric composition, variability corrupting pan-epoch colors, and suspected typographical errors in published literature values.  \nFigure 9. Various color-color diagrams for the 20-pc census. See text for details.  \n<!-- image -->  \n- · HD 1237 B (0016 -7951): For its spectral type of M4, this object has GBP -G and GBP -GRP colors much bluer than expected, while its G -GRP color is much redder than expected. No other separate photometry of the B component is given in Table 4. Given that the A component is eight magnitudes brighter in G than the B component and lies only 4 . \'\' 0 away, we suspect a problem with the measured photometry of B that is not adequately reflected in its quoted uncertainties.\n- · EGGR 246 (0041 -2221): This is an oddly blue white dwarf in colors that use W1 or W2 magnitudes. The object is also blue relative to other white dwarfs in J -H color, though normal in G -J . This carbonbearing object has a peculiar spectral type, DQpec9.3, and is believed to have a mixed hydrogen-helium atmosphere. The known infrared flux deficit is thought to be caused by absorption by H2 via collisions with neutral He (Giammichele et al. 2012; Bergeron et al. 1994, 2022).\n- · LP 941-19 (0213 -3345): Although this DA4.5 white dwarf has contaminated WISE photometry, it falls in an odd position on plots based only on Gaia photometry. Specifically, at its value of MG it falls ∼ 0.5 mag blueward of the white dwarf locus in GBP -G and ∼ 0.3 mag redward in G -GRP . There is very little literature on this source, and our spectrum of it (Section A.2) is the first published. It is not yet clear if this spectrum differs markedly enough from other DA white dwarfs to account for the color discrepancies or whether the Gaia magnitudes themselves are at fault.\n- · HD 21209 A (0323 -4959): The only oddity with this K dwarf is its anomalously blue W1 -W2 color. The value in Table 4, which is from the WISE All-Sky Source Catalog, is W1 -W2 = -0 . 17 ± 0 . 06 mag. Although this is the preferred WISE catalog for sources of this brightness (W1 = 5.56 mag), the AllWISE Source Catalog gives a very similar color of W1 -W2 = = -0 . 15 ± 0 . 13 mag. This color may be due to the slightly subsolar metallicity of the object ([Fe/H]= -0 . 44 ± 0 . 19, Soto & Jenkins 2018; -0 . 41 ± 0 . 04, Sousa et al. 2008; -0 . 39 ± 0 . 02, Tsantaki et al. 2013).\n- · HD 23189 (0348+6840): This early-K dwarf is underluminous for its type at MG , MJ , MH , MKs , and MW 2. When colors formed from Gaia-based magnitudes are compared to the mean colors of objects of the same type, it appears normal, whereas the W1 -W2 color is slightly bluer than normal. We suspect that the Gray et al. (2003) type of K2 V is a typographical error, as independent assessments of the type from spectra, colors, and luminosity considerations (Adams et al. 1935; Bidelman 1985; Mermilliod 1987; Stassun et al. 2019) suggest a spectral type closer to K7.\n- · 2MASS J05053461+4648017 (0505+4648): This is an M8 dwarf (see Section A.2) with a Gaia DR3 parallax value of 56 . 84 ± 0 . 60 mas. The absolute values calculated with this parallax are similar to those of other known M8 dwarfs in the census, and a previous parallax of 69 . 5 ± 4 . 7 mas (Dittmann et al. 2014) also places it within 20 pc. Its location on color-magnitude diagrams such as MG vs. G -W2 along with its high DR3 RUWEvalue of 4.823 indicate possible unresolved binarity.\n- · DENIS J071807.3 -350220 (0718 -3502): On the GBP -GRP and G -J vs. various absolute magnitude diagrams, this object appears to be ∼ 0.7 mag above the locus of other objects of the same color. This is, therefore, likely a near equal-magnitude double. This object is also flagged as a possible binary in the Apps Catalog, again based on its position in color-magnitude diagrams.\n- · SCR J0818 -3110 (0818 -3110): This DZ white dwarf is an outlier on the G -J vs. W1 -W2 diagram and, in fact, any diagram involving W1 -W2 color. This issue has been hinted at previously in Figure 4 of Kawka et al. (2021), which shows that the best model fit to existing spectra and photometry fails to match the observed W1 -W2 color. Although the effect is known, its reason has apparently not yet been established and may be caused by variability or missing opacity sources in the models.\n- · UPM J0901 -6526 (0901 -6526): This K5 dwarf is an outlier on all plots showing spectral type but appears normal on color-color and color-magnitude plots. We suggest that the published type of K5 (Riaz et al. 2006) results from a transcription error in the data for this star and that the actual spectral type is closer to M5.\n- · APMPM J1251 -2121 (1250 -2121): This M6/6.5 dwarf has a Gaia DR3 parallax of 56 . 79 ± 0 . 19 mas. Its apparent magnitudes fall within the range expected for an M6 dwarf within the 20-pc volume, and a previous parallax measurement of 57 . 7 ± 1 . 7 mas (Winters et al. 2015) is in agreement with the Gaia one. On the MG vs. GBP -GRP diagram, it falls ∼ 0.5 mag above the main sequence, and on the MG vs. G -W2 diagram, it falls ∼ 0.7 mag above. This and the Gaia DR3 RUWE value of 2.888 suggest unresolved binarity.\n- · HD 113194 (1302 -2647): Although this K5 dwarf has a Gaia DR3 parallax with a relatively large uncertainty (56 . 94 ± 0 . 19 mas), the earlier Hipparcos parallax (56 . 87 ± 1 . 11 mas) is in agreement. This object also has a high Gaia DR3 RUWE value, is listed as a high-LUWE object (see Section 5.2), and shows acceleration (see Section 4.1.4), in addition to falling ∼ 0.5 mag above the main sequence on the G vs. GBP -GRP diagram. This object is almost certainly an unresolved binary. It is also considered to be binary in the Apps  \nCatalog, based on its position on color-magnitude diagrams. The Gaia DR3 main catalog reports a radial velocity of -17 . 56 ± 7 . 24 km s -1 using seventeen observations over 920 days, along with an amplitude of radial velocity variations of 60.25 km s -1 , further supporting the hypothesis of binarity. The P-value for radial velocity constancy ( rv\\_chisq\\_pvalue ) is also 0.0.  \n- · 2MASSW J1421314+182740 (1421+1827): This M9.5 dwarf has WISE photometry contaminated by a background source, but it appears unusual in nonWISE colors as well. Specifically, it has oddly blue GBP -G and GBP -GRP colors compared to other objects of similar MG . However, it appears normal for its MG in G -GRP , G -J , G -H , and J -Ks . This may simply indicate an issue with the GBP magnitude that the formal uncertainty fails to adequately capture.\n- · LP 222-65 (1516+3910): This mid-M dwarf lies consistently ∼ 0.6 mag above the main sequence relative to objects of the same color and spectral type on colormagnitude diagrams. This is an isolated object with no obvious problems with its photometry, so we believe this is an unresolved near-equal magnitude binary.\n- · UCAC4 554-051865 (1518+2036): This mid-M dwarf is ∼ 0.6 mag more luminous than objects of similar color on the MG vs. GBP -GRP and MG vs. GBP -G diagrams and has a large RUWE and LUWE value (see Section 5.2). It is likely an unresolved binary.\n- · L 339-19 (1640 -4559): This M3 dwarf shows anomalously red G -W3 and J -W3 colors for its absolute magnitude, and even more anomalously red G -W4and J -W4 colors. A more careful look at the WISE images shows that the W3 detection is likely real, but the W4 detection likely is not. The W3 photometry from AllWISE (reported in Table 4) is 7.10 ± 0.05 mag and that from WISE All-Sky is 6.57 ± 0.04 mag. In G -W3 and using the AllWISE value, the object lies 1.0 mag redward of objects of the same absolute G magnitude; using the WISE All-Sky value shows the object to lie 0.5 mag redward. Archival Spitzer/IRAC and Spitzer/MIPS photometry of this object exists in the GLIMPSE I Spring \'07 Catalog (Benjamin et al. 2003) and MIPSGAL Archive (Carey et al. 2009) at IRSA: ch1 = 7.830 ± 0.038, ch2 = 7.781 ± 0.045, ch3 = 7.724 ± 0.037 (5.8 µ m), ch4 = 7.705 ± 0.026 (8.0 µ m), and [24 µ m] = 7.14 ± 0.24 mag. Running these new data points and the tabulated Table 4 photometry through the Virtual Observatory Spectral energy distribution Analyzer (VOSA 31 ; Bayo et al. 2008) suggests not only that the W4 magnitude is in error but that the W3 magnitude is spuriously bright relative to  \nthe bracketing IRAC and MIPS data points. The spectral energy distribution is otherwise typical of that of an M3 dwarf. We therefore conclude that there is no infrared excess in this object.  \n- · UCAC4 317-104829 (1706 -2643): This DAH white dwarf is normal in Gaia-only colors, colors formed using Gaia minus near-infrared magnitudes, and colors formed from J , H , and Ks magnitudes. It is, however, oddly blue in W1 -W2. We assume that this anomalous color may be intrinsic to the star and a result of its strong magnetic field, although it should be cautioned that this white dwarf is located against a busy region of the Galactic Plane and may suffer from contamination in its WISE photometry.\n- · DENIS-P J1733423 -165449 (1733 -1654): This L1 dwarf has WISE photometry that is contaminated by background sources, but it also shows unusual colors in Gaia-only measurements. Gaia DR3 lists two other point sources within 2 . \'\' 1 of this object, so its Gaia photometry may be adversely affected in a way that the formal uncertainties fail to capture.\n- · LSPM J1733+1655 (1733+1655): The Gaia DR3 parallax of 60 . 91 ± 0 . 48 mas has a relatively large uncertainly for its magnitude and is in disagreement with an earlier published value of 85 . 40 ± 3 . 30 mas by Dittmann et al. (2014). This mid-M dwarf is more luminous than objects of similar color by ∼ 1.6 mag on the G vs. GBP -G , GBP -GRP , G -GRP , G -J , and G -W2 diagrams, if the Gaia DR3 parallax is used. This overluminosity decreases to ∼ 0.9 mag if the Dittmann et al. (2014) parallax is used instead. This is a high-RUWE/LUWE object as well (Section 5.2), and so is likely an unresolved multiple system with problematic Gaia astrometry. Clark et al. (2022) identify a candidate companion at separation 0 . \'\' 14 and position angle 101 · at epoch 2017.3 and again at separation 0 . \'\' 36 and position angle 63 · at epoch 2019.7. C. Gelino also finds a single epoch of Keck/NIRC2 data in the Keck Observatory Archive for LSPM J1733+1655. These are Br γ and J -continuum observations taken on 2015 Jul 10 UT (PI: Hansen; Program ID: U050N2), from which we measure a separation of 0 . \'\' 11 at position angle 248 · . If we assume all three of these measurements refer to the same star and it is a stationary background object, we derive motions of LSPM J1733+1655 of -0 . 100 \'\' yr -1 in RA and -0 . 051 \'\' yr -1 in Dec, which can be compared to the measured Gaia DR3 values of -0 . 135 \'\' yr -1 in RA and -0 . 130 \'\' yr -1 in Dec. The derived magnitude and direction of motion lead us to conclude that the background star hypothesis is sound.\n- · LP 388-55 A (1735+2634): This late-M dwarf is anomalously red, by 0.25 mag, in G -GRP color but looks normal in the GBP -G color compared to objects  \nof similar MG magnitude. Curiously, all Gaia-based absolute magnitudes ( MG , MGBP , and MGRP ) are consistent with the reported spectral type. The B component is an early-L that is not directly imaged by Gaia but may nonetheless be subtly affecting the Gaia magnitudes of the A component.  \n- · LP 44-334 A (1840+7240): This primary in a M6.5 dwarf system has a Gaia DR3 parallax (52 . 78 ± 0 . 09 mas) with a relatively large uncertainty for its magnitude, but this value compares favorably to the earlier published value of 59 . 3 ± 2 . 2 mas by Lépine et al. (2009). The GBP -G color is too blue for its MG value, the GBP -GRP color is normal, and the G -GRP color is too red. These issues are likely caused by the nearness of the B component, only 0 . \'\' 8 away, which is likely corrupting the photometry of the A component.\n- · LP 867-15 (1842 -2328): The colors for this M0 dwarf are more consistent with an M4 dwarf than with an M0. Pending spectroscopic verification, we assume that this object has been misclassified.\n- · SCR J2012 -5956 (2012 -5956): This object, a DC9.9 white dwarf, falls below the white dwarf locus for most colors. It is very blue relative to other white dwarfs in J -Ks , J -H , and H -Ks but looks like other white dwarfs in colors made with Gaia-only magnitudes. It is somewhat blue in G -J , G -H , and G -Ks colors. As with EGGR 246 above, the infrared flux deficit is believed to be caused by H2-He collision-induced absorption (Giammichele et al. 2012).\n- · LEHPM 2-783 (2019 -5816): This M6.5 dwarf is overluminous in all Gaia-based colors. (Many other colors are nearly degenerate with absolute magnitude or type in this spectral type range.) On both the GBP -GRP vs. MH and the G -W2 vs. MW 2 plots, the overluminosity is ∼ 0.7 mag. Ujjwal et al. (2020) mark this as a possible member of the β Pic Moving Group, and Riaz et al. (2006) note that it is a strong X-ray emitter with strong H α emission.\n- · LP 12-90 (2322+7847): This mid-M dwarf lies above the main sequence by ∼ 0.75 mag on the MG vs. GBP -G plot. On many other plots of absolute magnitude vs. color, it lies similarly above (and redward of) the main sequence. This could be another unresolved binary - if confirmed, this would make its system with HD 220140 AB a quadruple - but the primary in this system is a young, naked T Tauri star (Makarov et al. 2007), meaning that its position may be solely due to its youth.\n- · ZZ Psc (2328+0514): This white dwarf is anomalously red in colors involving WISE magnitudes - so much so that it falls far from the white dwarf sequence itself. On a plot of MG vs. W1 -W2, for example, it  \nlies substantially redward of both the white dwarf locus and the main sequence. This object, also known as G 29-38, is known to have a debris disk around it, the first evidence of which was uncovered by Zuckerman & Becklin (1987). For an update on this object, see Cunningham et al. (2022).', '6. MASSES FROM ESTIMATION': 'Only a small fraction of objects within the 20-pc census has masses measurable by methods 1 or 2 described in the introduction of Section 4. For the rest, we must rely on methods 3 and 4 of that section, which depend on comparison to empirical trends or to theoretical models.  \nIn the first subsection below, we discuss mass measurements for objects not on the main sequence - namely, white dwarfs, giants/subgiants, and brown dwarfs. In the second subsection, we summarize mass estimation for main sequence stars. In the third subsection, we discuss other complications - youth, subsolar metallicity, and formation scenario - that may need to be considered when assigning accurate mass estimates for special objects.', '6.1.1. White dwarfs': "Masses have been measured via one of the methods described in Section 4 for a handful of white dwarfs in the 20-pc census, but these represent the end-state masses of the stellar remnants and are not suitable for analysis of the initial mass function. Rather, what is needed are the initial masses before evolution off the main sequence. Techniques have been established that use the final mass of the remnant to estimate the initial mass of the progenitor.  \nFor white dwarfs lacking a direct mass measurement, one can estimate the final mass of the white dwarf using one of the following two semi-empirical methods. The first is to use spectroscopic observations of the depth and width of the hydrogen Balmer, He I, or He II lines to establish, after comparison to atmospheric models, the log( g ) and T eff for each object. Further comparison of these two parameters to cooling models provides the remnant mass (e.g., Tremblay et al. 2011; Genest-Beaulieu & Bergeron 2019; Bergeron et al. 1992; Finley et al. 1997). Whereas this first method is applicable only to DA (hydrogen atmosphere) or DB (helium atmosphere) white dwarfs, an alternate method can be used both for objects lacking hydrogen lines as well as for objects lacking spectroscopic observations. In this second method, masses can still be estimated if an accurate parallax has been measured. Here, absolute fluxes across as wide a swath of wavelength space as possible are compared to model atmospheres to provide log( g ) and R , from which the mass can be derived from equation 11 (e.g., Bergeron et al. 2019; Tremblay et al. 2019; Giammichele et al. 2012; Bergeron et al. 2001; Koester et al. 1979).  \nThe next step is to convert this final mass into an initial mass using an initial-to-final mass relation (e.g., Weidemann 2000). The empirical form of this relationship has been es-  \ntablished using white dwarfs that are members of open clusters of known age. As described above, spectroscopic observations of the Balmer lines in these stars can be compared to atmospheric models to derive log( g ) and T eff for each object. A comparison of these parameters to cooling models provides both the remnant mass as well as the cooling time since the object left the tip of the asymptotic giant branch. The known cluster age minus this cooling time gives the main sequence lifetime of the object, which can then be related back to an initial mass using theoretical evolutionary isochrones. This same technique can also be applied to white dwarfs in globular clusters. Due to their much older ages, these clusters can provide white dwarfs of lower final mass than those available in young open clusters. Because these globular clusters are much more distant, their white dwarfs are faint and more difficult to study, so old low-mass white dwarfs are still not well represented by cluster methods.  \nThis lack of low-mass examples can be partly mitigated by the use of old, wide binaries for which the second component can be age dated and the separation between components is large enough that no mass transfer has occurred during the system's evolution. Examples are wide subgiant + white dwarf binaries in which the system can be dated from its more recently evolved member (Barrientos & Chanamé 2021), wide F/G/K dwarf + white dwarf binaries in which the age of the main sequence star can be estimated from activity diagnostics (Catalán et al. 2008, Zhao et al. 2012), and white dwarf pairs in which comparison of the higher-mass white dwarf to known cluster white dwarfs can provide an age for the binary, and the difference in the cooling times for the white dwarf pair gives the main sequence lifetime of the lower-mass white dwarf (Andrews et al. 2015). Using the results of these methods, the trend of final mass with initial mass can be fit. As can be seen from figure 9 of Bar-  \nentos & Chanamé (2021), the relation shows considerable scatter at lower masses, as the age dating methods for individual systems are generally less robust than those from clusters. The relations we adopt here are the cluster-based tripartite parameterization found in equations 4-6 of Cummings et al. (2018) and the quadripartite parameterization found in table 1 of El-Badry et al. (2018), The former relation is applicable to white dwarfs with 0 . 56 < M final < 1 . 24 M ⊙ (0 . 83 < M initial < 7 . 20 M ⊙ ), and the latter relation, which is based on nearby white dwarfs with accurate Gaia parallaxes, is applicable to white dwarfs with 0 . 50 < M final < 1 . 37 M ⊙ (0 . 95 < M initial < 8 . 00 M ⊙ ). We further note that neither the cluster nor field methods have yet extended the initial-to-final mass relation below final masses of 0 . 50 M ⊙ . (As discussed further below, white dwarfs with final masses below 0 . 45 M ⊙ require binary interactions, as a single progenitor would imply an age older than the Universe; Marsh et al. 1995).  \nSpecifically, we apply the following methodology to assign final masses to white dwarfs in the 20-pc census. First, we use directly measured masses, whenever such measurements are available. For others, we use final mass estimates that are based on accurate parallaxes, high S/N spectra, and/or broad-wavelength data spanning the white dwarf's spectral energy distribution. For all other objects, we resort to the Gaia-centric estimates of Gentile Fusillo et al. (2021) and Gentile Fusillo et al. (2019). These estimates use only a small fraction of the white dwarf's spectral energy distribution spanning the Gaia optical bandpasses - and thus lead to separate solutions for hydrogen- vs. helium-atmosphere objects. When our own follow-up has determined the spectral type of the object, we use this information to break the degeneracy; otherwise, a hydrogen-atmosphere object is assumed, as noted in Table 10.  \nTable 10 . Mass Measurements and Estimates for White Dwarfs in the 20-pc Census", '6.1.2. Giants and subgiants': 'There are a number of objects in the 20-parsec census that have evolved off the main sequence but have not yet become white dwarfs. Table 11 includes all objects in Table 4 that have a luminosity class more luminous than V and/or fall in a locus on the absolute magnitude vs. color diagrams that identifies them as post-main sequence stars.  \nSeveral of these have direct mass measurements either from orbital dynamics or asteroseismology. The rest have had their masses estimated from other methods, primarily via comparison of their placement on the HR diagram in relation to modeled evolutionary tracks or via fits of their spectra to atmospheric models. For some objects with IV-V or IV luminosity classes, other published spectral types indicate a V luminosity class or their placement on the HR diagram suggests a main sequence star. This is reflected in the mass estimates given in Table 11.  \nTable 11 . Giants and Subgiants in the 20-pc Census  \nTable 11 continued', '6.1.3. Brown dwarfs': 'Brown dwarfs follow no mass-luminosity relation because they constantly cool over time. If the age of the brown dwarf is known, this can be used to estimate the mass from evolutionary models, but age is a difficult parameter to measure for non-youthful disk objects. We therefore must resort to simulations to tease out information regarding the mass function. In Kirkpatrick et al. (2019, 2021a), we took the empirical distribution of brown dwarf effective temperatures and compared that to various predicted temperature distributions modeled by taking the shape of the brown dwarf mass function, the value of its low-mass cutoff, and the underlying evolutionary model suite as free parameters. For the analysis of this paper, we will employ those same methods, using an updated suite of predictions by Raghu et al. (submitted).  \nHere, we compare the Kirkpatrick et al. (2021a) accounting of all 525 known 20-pc L, T, and Y dwarfs to that given  \nin Table 4. Additions and subtractions to this tally are listed in Table 12. We find that eight objects have fallen out of the 20-pc sample, all because of new parallax measurements or revised distance estimates that place them outside of 20 pc. On the other hand, sixty-five objects are newly added. These additions include thirty-eight new discoveries (thirtyseven by the Backyard Worlds citizen science group, four of which are new companions), nine new companions recently announced in the literature, one new companion announced here but found in Gaia, three new published parallaxes with  \nd < 20 pc, twelve previously overlooked companions, and two previously overlooked objects (DENIS J065219.7-253450, presumably due to a transcription error, and SSSPM J14442019, whose subdwarf type had earlier been updated from late-M to early-L). To facilitate analysis on the revised T eff distribution, we have listed in Table 12 the estimated temperatures of each of the additions and subtractions. Further analysis can be found in Section 7.  \nTable 12 . Additions to and Subtractions from the 20-pc L, T, and Y Census of Kirkpatrick et al. (2021a)  \nTable 12 continued', '6.2. Other complications': "6.2.1. Youth  \nWill the estimation of masses for young objects be biased if those estimates use a relation based on much older stars? Evolutionary models suggest that below a mass of ∼ 0 . 4 M ⊙ , the contraction of a star down to the main sequence follows a Hayashi track along which the star's effective temperature remains approximately fixed (section 16.2.5 of Stahler & Palla 2004). If a temperature-based metric is used for estimating the masses of such stars, then such estimates will be accurate. At higher masses, however, the descent along the Hayashi track will be interrupted when a radiative zone develops. The star then moves via a Henyey track along which the temperature slowly increases until the star reaches the main sequence. For stars with masses above ∼ 0 . 4 M ⊙ , this evolution to the main sequence occurs within the first 100 Myr.  \nThis means that objects in the 20-pc census that have masses above ∼ 0 . 4 M ⊙ and ages less than 100 Myr should have their mass estimates more carefully considered. The  \nMontreal Open Clusters and Associations database (https: //mocadb.ca/, Gagné et al., in prep., Gagné et al. 2018) is a compilation of known stellar associations, stellar streams, moving groups, and open clusters within 500 pc of the Sun. A search of this database on 2023 May 18 for objects within 20 pc of the Sun and likely belonging to one of these young groups yielded 217 systems. The only objects in this list with ages below 100 Myr are those believed to be members 37 of the β Pic Moving Group ( ∼ 26 Myr), the Columba Association ( ∼ 42 Myr), the Argus Association ( ∼ 45 Myr), the Carina Assocation ( ∼ 45 Myr), and the Octans-Near Association ( ∼ 55 Myr) 38 . These are listed in Table 13.  \nBecause a mass of 0 . 4 M ⊙ corresponds to a spectral type of M2.5-M3 (table 7 of Mann et al. 2019), we can use spectral type to identify which of the young 20-pc objects are the ones whose mass estimations may need special handling.  \nTable 13. 20-pc Members of Young Associations and Moving Groups with Ages < 100 Myr  \nThe only objects in Table 13 with spectral types earlier than this are BD -21 1074 ABC, V2689 Ori, β Pic, α Cir AB, HD 182488 A, AU Mic, and HD 220140 A. Three of these are early M dwarfs for which the brief jog along the Henyey track before reaching the main sequence covers such a small range in temperature that their mass estimates should not be unduly affected.  \nThis leaves only four individual objects to consider, and two of these have dynamical mass measurements already. For β Pic, Lacour et al. (2021) used astrometry of the exoplanet system to derive the mass of the host star (1.75 ± 0.03 M ⊙ ) using only a uniform mass prior between 1.4 and 2.0 M ⊙ on β Pic itself. HD 182488 A has a loosely constrained dynamical mass measurement from Brandt et al. (2019) of 0 . 94 + 0 . 17 -0 . 27 M ⊙ . This leaves only two young systems with possibly skewed mass estimates, and this represents such a small percentage of 20-pc stars with types earlier than M0 ( < 1%) that no bias will be imparted on the overall derived mass distribution.  \nWe acknowledge that our understanding of young moving groups near the Sun is still evolving. Our Sun is currently moving through three groups - the β Pic Moving Group, the AB Dor Moving Group, and the recently identified (but older) Oceanus Group (Gagné et al. 2023) - but it remains unlikely that many new early-M and hotter dwarfs within 20 pc will be associated with any newly recognized groups. Such young objects would have already revealed themselves through, for example, high chromospheric activity.  \nYoung brown dwarfs , on the other hand, require their own special handling. For brown dwarfs, we deduce the form of the mass function via the empirical temperature distribu-  \ntion. It has been well established, however, that young brown dwarfs follow a different spectral type (or color) to T eff relation than their older counterparts (Faherty et al. 2016). Corrections to the temperature estimates for these objects were already established for the brown dwarf portion of the 20pc census in Kirkpatrick et al. (2021a), and none of the new brown dwarfs discussed in Table 12 are known to be youthful themselves. Therefore, no additional work is required here.", '6.2.2. Non-solar metallicity': 'Objects with non-solar metallicity raise two concerns. The first is that metal-poor objects may belong to the Galactic halo population and could skew our calculation of the nearby mass function, which concentrates on the Galactic disk. The second is that these objects, even if true disk members, may be sufficiently metal poor that standard solar-metallicity relations will not adequately predict their masses. Are either of these concerns justified?  \nA number of objects in Table 4 have spectroscopic classifications indicating subsolar metallicity. For objects earlier than early-M, these classifications can generally be identified via the iron index, "Fe#", which attempts to encode the abundance of metals relative to hydrogen in the spectrum if the spectrum does not match the standards of solar-metallicity (Gray & Corbally 2009). Underabundances are encoded as negative numbers. For objects of spectral type late-K and later, metal poor spectral types (Gizis 1997; Lépine et al. 2003; Kirkpatrick 2005; Lépine et al. 2007; Burgasser et al. 2007c; Zhang et al. 2017) are usually denoted with prefixes of sd (subdwarf), esd (extreme subdwarf), or usd (ultra subdwarf). Table 14 lists all objects in the 20-pc census that have one of these low-metallicity classifications.  \nTable 14 . 20-pc Objects with Low-metallicity and/or Halo Kinematics  \nTable 14 continued  \nTable 14 (continued)  \nTable 14 continued  \nTable 14 (continued)  \nReferences -The references for radial velocity are - (1) Gaia Collaboration et al. 2018, (2) Fouqué et al. 2018, (3) Gaia Collaboration et al. 2022, (4) Holmberg et al. 2007, (5) Pourbaix et al. 2004, (6) Maldonado et al. 2010, (7) Abazajian et al. 2009, (8) Gizis 1997.  \nothers most likely belong to the thin or thick disk populations.  \nTo answer the first concern, we use the sky positions, parallaxes, and proper motions in Table 4 along with published radial velocities in Table 14 to calculate the U , V , W space velocities with respect to the Local Standard of Rest (LSR). We also calculate the U , V , W values for all objects in Table 4 with Gaia-based radial velocity measurements to see if any objects lacking low-metallicity spectral classifications are found to be halo members merely from their kinematics 39 . Figure 10 shows the Toomre diagram for both sets of objects. Also shown for comparison are stars having radial velocity measurements in Gaia DR2 and lying within 100 pc of the Sun, color coded as thin disk ( V tot ≤ 85 km s -1 ), thick disk (85 < V tot ≤ 180 km s -1 ), or halo ( V tot > 180 km s -1 ) in accordance with the kinematic criteria of Nissen (2004). This comparison demonstrates that only six objects - LP 651-7, Ross 578, HD 25329, Kapteyn\'s Star, HD 103095, and Ross 769 - appear to belong to the kinematic halo population. All  \nAs stated in Section 3.5, Gaia contains radial velocities only for those objects having GRVS ≲ 14 mag, which omits many of the M dwarfs and all of the L, T, and Y dwarfs within 20 pc. For these objects, we leverage spectroscopic indications of low metallicity to build a list of potential halo members, then we scour the literature for other published radial velocities. These objects are also listed in Table 14. Many of these lack any radial velocity measurements, so assumed values from -200 to +200 km s -1 , in increments of 50 km s -1 , were used to calculate a range of possible U , V , W velocities. These results, shown in Figure 11, suggest that only two of these colder objects - SSSPM J1444 -2019 and WISEA J153429.75 -104303.3 (aka "The Accident") - are likely to be true halo members.  \nFigures 10 and 11 taken together suggest that only eight objects (all of them believed to be single) out of 3,589 total in the 20-pc census, or 0 . 22%, are halo interlopers. Although this is slightly higher than the percentage of 0 . 15% used in Table A of Bensby et al. (2014) based on F and G stars alone, it nevertheless confirms that contamination by halo objects in the 20-pc census is extremely small. Although these objects  \nFigure 10. Toomre diagram of UVW space motions corrected to the Local Standard of Rest (LSR) for 74,066 Gaia DR2 stars within 100 pc of the Sun and having parallax errors < 10% (Kirkpatrick et al. 2021b). Thin disk (light grey dots), thick disk (medium grey crosses), and halo (dark grey pluses) objects are marked, with halo stars falling outside the outer dashed circle (red) and thin disk objects falling inside the inner dashed circle (navy). Objects with measured radial velocities in Table 4 or Table 14 are shown in navy if lying in the thin disk velocity zone, yellow for the thick disk zone, and red for the halo zone. The six halo members are highlighted with black labels.  \n<!-- image -->  \nwill still be included in our mass function, any systematic offset imprinted upon their mass estimates can be ignored in subsequent analyses.  \nThe second concern is difficult to address, as very few lowmetallicity objects have had their masses measured via direct methods. The coldest subdwarfs, for instance, have a multiplicity fraction of only ∼ 1% (González-Payo et al. 2021); therefore, few such objects exist for dynamical analyses (e.g., Rebassa-Mansergas et al. 2019). Single subdwarfs are obvious targets for lensing-based mass measurements, as their high velocities increase the likelihood of "encounters" with background objects, but accurate whole-sky astrometry is just now advancing to the stage at which such measurements can be predicted and planned for (e.g. Sahu et al. 2020). So, to address this concern, we instead note that only forty-two  \nFigure 11. Toomre diagram of the 100-pc sample from Figure 10, now overplotted with the seven objects (various colors and symbols) from Table 14 that lack radial velocity measurements. For these, results are shown for nine assumed radial velocities ranging from -200 to +200 km s -1 , in steps of 50 km s -1 . As in Figure 10, the demarcation of the thin disk, thick disk, and halo populations are shown by the dashed circles in red and navy.  \n<!-- image -->  \nlow-metallicity systems are known within the 20-pc census 40 (Table 14), which represents only 1 . 5% of the total. Thus, if small biases are present in converting a subdwarf\'s spectral type, colors, or absolute magnitudes to masses, the bias in the overall 20-pc mass distribution will be negligible.  \nThe above logic on the scarcity of objects also holds for systems with higher metallicity than the Sun. This set of objects has a much smaller range in metallicity than the metalpoor objects above, and there are just a handful of examples. Only the higher-mass stars ι Hor AB (Fe+0.3), ν Phe (Fe+0.4), HD 176051 AB (slightly metal strong), and HD 207129 (Fe+0.4) have spectroscopic classifications that fall into this class. Another object, 14 Her, has a supersolar metallicity ([Fe/H] ≈ 0 . 4; Rosenthal et al. 2021) although its listed spectral type in Table 4 gives no indication of this. Curiously, even though members of the Hyades Cluster have metallicities that are slightly supersolar ([Fe/H] = 0 . 14 ± 0 . 05;  \nPerryman et al. 1998) and lie, on average, only 47.0 ± 0.2 pc from the Sun (Lodieu et al. 2019), there are no confirmed Hyads within the 20-pc volume (Gaia Collaboration et al. 2020; Schneider et al. 2022).', '6.2.3. Formation process': 'Because we are interested in objects formed via the star formation process, we need a criterion to distinguish objects that may have formed via alternative formation mechanisms at the lowest masses. When brown dwarfs were first theorized (Kumar 1963; Hayashi & Nakano 1963), they were regarded as direct products of the star formation process - ones that had insufficient mass to sustain prolonged thermonuclear fusion in their cores - and as such represented the lowermass extension of hydrogen-burning stars themselves. These could be contrasted with another low-mass formation product, planets, that were believed to be formed via a secondary process - from a protoplanetary disk created around a newly formed protostar or brown dwarf. In the early 1960s, there were no known examples of brown dwarfs, and our own Solar System provided the only known examples of planets.  \nAs brown dwarf and exoplanet discoveries began in earnest (see reviews by Kirkpatrick 2005; Winn & Fabrycky 2015), it became clear that nature produces some low-mass products that are difficult to classify as either brown dwarf or exoplanet (e.g., 2MASSWJ 1207334 -393254b, Chauvin et al. 2004). The earlier definition based on formation was cumbersome to use in practice; unless an object was still in its infancy, its exact formation process would be difficult, if not impossible, to ascertain from observations. As an alternative, Burrows et al. (1997) proposed another theoretically based definition. This alternative uses mass to distinguish between a brown dwarf and an (exo)planet, the dividing line being the somewhat arbitrarily chosen deuterium burning limit, which is ∼ 13 MJup for solar metallicity. Somewhat surprisingly, this definition was thereafter widely (though not universally) adopted, in no small part because lower-mass discoveries that earlier would have been called "brown dwarfs" could now be referred to by a more attention-grabbing label of "exoplanet."  \nThis alternative definition, however, came three and a half decades after the original brown dwarf definition, and the concept of planets having being born from a protoplanetary disk (the "nebular hypothesis") had been in the astronomical lexicon for over two centuries (Kant 1755; Laplace 1796). Thus, labeling an object below 13 MJup as a planet often leads to confusion, as some readers - and even researchers - unwittingly apply both definitions in tandem. That is, they assume that a so-named "planet" (by the new, mass-based definition) must have formed via a protoplanetary disk (by the former, formation-based definition). It is difficult to divorce the term "planet" from its formation scenario.  \nIn this paper, one of our goals is to define, or place limits on, the low-mass terminus of star formation. If we were to use the newer definition to include/exclude objects for the mass function analysis, our results would return a terminus of 13 MJup , which merely reflects the dividing line chosen by the arbitrary definition. We must, therefore, more carefully  \nconsider whether the lowest mass objects in the 20-pc census should be counted as star-formation products or planetaryformation products.  \nAs stated earlier, this definition also lacks easy observational verification. Nonetheless, some methods have been proposed to distinguish formation mechanisms. Öberg et al. (2011) postulated that the carbon to oxygen ratio could be used as one tracer. Planets that formed close to a star would have a solar-like C / O value, like brown dwarfs formed via gravitational collapse, whereas planets formed via accretion of ices beyond the water snowline would have a supersolar C / O value. Those authors acknowledged, however, that measuring an accurate value of C / O is fraught with difficulties (even within our own Solar System), and Calamari et al. (2022) made a similar conclusion based on their analysis of the spectrum of the brown dwarf Gliese 229B. Mollière et al. (2022) show that this simplified picture of the C / O ratio is somewhat more complicated when disk chemical evolution and pebble accretion are taken into account, as well.  \nSimilarly, Morley et al. (2019) showed that the deuterium to hydrogen ratio could be used to distinguish between planets with solar D / H values like Jupiter and Saturn, that formed directly from accretion of gas in the protostellar nebula, and planets with enhanced D / H values like Neptune and Uranus, that presumably formed from accretion of ices. Both C / O and D / H thus have limitations: some objects formed via a protoplanetary disk have values indistinguishable from those of objects born via star formation.  \nAnother promising avenue is the overall metallicity. The giant planets of our Solar System have metal enhancements well above solar values (Wong et al. 2004, Fletcher et al. 2009), and exoplanets are preferentially found around metalrich host stars (Fischer & Valenti 2005, Wang & Fischer 2015). These facts led Fortney et al. (2008) to propose metallicity-based diagnostics that could distinguish between formation scenarios. Specifically for objects with T eff < 1400 K , a strong 4.5 µ m CO absorption band along with enhanced H - and K -band fluxes (from a relative lack of collision-induced absorption by H 2) are proposed as fingerprints of planet-like formation. However, these diagnostics are likely only useful when comparing populations of objects and not when establishing the formation pathway of individual objects. Metal enrichment is not unique to planet formation, as a collapsing metal-rich cloud can also produce low-mass objects.  \nBowler et al. (2023) note that the orientation between the spin axis of the star and the orbital plane of the companion shows promise as another marker of formation, as starlike formation shows a wide range of orientations, whereas planet-like formation prefers values near 90 · . This is, however, another marker that can distinguish between populations but cannot be used on an individual object basis.  \nSchlaufman (2018) demonstrates that companions above ∼ 10 MJup lack the tendency to fall primarily around metalrich hosts that companions below ∼ 4 MJup exhibit, which is taken as evidence of core accretion in the lower-mass set. Hoch et al. (2023) likewise find a tentative difference in the  \ntrend of C / O values at ∼ 4 MJup , which is taken as further evidence that those objects are primarily formed via core accretion, although, as stated above, C / O ratios can be difficult to interpret. Similarly, Ribas & Miralda-Escudé (2007) find differing radial velocity distributions above and below M sin( i ) values of ∼ 4 MJup . Schlaufman (2018) states that planet-like formation appears to cease above ∼ 4 MJup , but not necessarily that star-like formation ceases below ∼ 10 MJup . There might still be a range in mass, below ∼ 4 MJup , where both processes contribute.  \nThe methods addressed above require data that are so far lacking for most exoplanets or can be used only in comparing populations. Instead, for this paper, we propose a simple scheme whose purpose is merely to exclude objects with a high likelihood of having been formed via a protoplanetary disk while including all others as possible products of star formation. For our scheme, we require at least three bodies in a system because the only parameters available for two-body systems - mass ratio, separation, etc. - can lead to ambiguities when trying to distinguish between formation scenarios.  \nAs an example, Bowler et al. (2020) have used twentyseven long-period companions labeled as giant planets and brown dwarfs to search for differences in parameters. They find that the population of brown dwarfs has an eccentricity distribution peaking in the range 0 . 6 < e < 0 . 9, whereas binaries with mass ratios significantly different from one have an eccentricity distribution peaking closer to e ≈ 0. These results indicate that the star formation process tends to create binaries with large eccentricity, and the protoplanetary process tends to form binaries with near-zero eccentricity. To reiterate a point from above, while such trends may be indicative of a population of objects, eccentricity alone cannot be used on an object-by-object basis to distinguish between formation scenarios. The same is true for mass ratio, as doing so can bias our list of potential companions to only the higher mass ones, which could impact our ability to determine star formation\'s low-mass cutoff. (Similarly, not se-  \ncting on mass ratio can bias our results in the opposite direction, a point we address further in the next section.)  \nIn triple (and higher-order) systems, however, we have other parameters available. Specifically, we note that the hierarchy of empirically observed triple star systems is such that the period of the outer component must be at least five times that of the inner pair (Tokovinin 2004). This is in good agreement with dynamical stability expectations for objects in circular orbits, and the period of the outer component must be even larger than five times the inner one when elliptical orbits are considered (Mardling & Aarseth 2001). Planets that have formed from a protoplanetary disk, on the other hand, can often arrange themselves in stable orbital configurations (e.g., in resonances with one another) that violate the above law. A multi-star system that formed via a collapsing cloud could, presumably, arrange itself in a similar manner if conditions were ideal, but such examples must be exceedingly rare. Therefore, we will use the ratios of orbital periods to identify "exoplanet" systems in the 20-pc census that most likely formed via a protoplanetary disk, and we retain all others for consideration as possible products of star formation.  \nTo this end, Table 15 lists all of the host objects from Table 4 that were labeled as having one or more confirmed exoplanets in the NASA Exoplanet Archive as of 2022 Sep 01 41 . For systems in which any pair of "exoplanets" violate the Pouter < 5 Pinner , we indicate the innermost pair that violates the rule and exclude all of the planets, thus including only the host star in later analysis. For all others, we have used the NASA Exoplanet Archive to compile their mass measurements. For objects identified only through radial velocity monitoring, we list the M sin( i ) values, since the inclination of the system is not known. For other objects - transiting systems, radial velocity systems with astrometric imaging, etc. - we list the actual measured masses. Incorporating these objects in to the stellar mass function analysis will be discussed further in Section 7.  \nTable 15 . 20-pc Objects Hosting "Planets"  \nTable 15 continued  \nTable 15 (continued)  \nTable 15 continued  \nTable 15 (continued)  \nTable 15 continued  \nTable 15 (continued)', '6.3. Objects on the main sequence': 'Main sequence objects with directly measured masses can be used to calibrate relations of mass vs. absolute magnitude or mass vs. spectral type. Studies have shown that the relation with the smallest intrinsic scatter for K and M dwarfs is the one using absolute K -band magnitude (Delfosse et al. 2000). The fact that the K -band relation shows the least scatter across the optical to near-infrared range is also predicted by model atmospheres, as this is the wavelength regime where competing physical effects modulated by metallicity variations largely cancel one another (Delfosse et al. 2000; Mann et al. 2019). More (and improved) dynamical mass measurements of binary stars 42 along with improved Gaia parallaxes have enabled Mann et al. (2019) to construct a mass vs. MKs relation that results in estimated masses with only 2-3% uncertainty. Specifically, Ks is used because 2MASS provides all-sky coverage at this band. We use the Mann et al. (2019) massMKs relation (their equation 2) over the range 5 . 0 ≤ MKs ≤ 11 . 0, roughly corresponding to spectral types from early-M to late-M 43 . These estimates and their propagated uncertainties are listed in columns "EstMassMKs" and "EstMassMKsErr" of Table 4.  \nFor other main sequence stars, we can use the methodology employed by Stassun et al. (2019). Using ∼ 20,000 (non-reddened) stars within 100 pc of the Sun with spectroscopically determined effective temperatures, they established a relation between T eff and GBP -GRP color. This is then mated with the results of Torres et al. (2010) that relate T eff to mass for stars with dynamically measured masses (Stassun et al. 2018b). This gives mass estimates with uncertainties of ∼ 6 . 4% (Stassun et al. 2019). We take mass estimates and their uncertainties directly from the revised TESS Input Catalog (TIC; Stassun et al. 2019) for stars in our Table 4. These values are listed in columns "EstMassTIC" and "EstMassTICErr". We note, however, that the Stassun et al. (2019) prescription for stars with T eff ≲ 4000K (see their appendix A.1 along with Muirhead et al. 2018) followed a different methodology. For these objects, masses were estimated using Ks magnitudes, Gaia DR2 parallaxes, and the Mann et al. (2019) mass-vs.MKs relation.  \nSome main sequence stars lack both Ks magnitudes and an entry in the TESS Input Catalog. For these, we resort to two other estimation methods. The first is the mass vs. MG relation. Chontos et al. (2021) took a list of well-studied late-K and M dwarfs (tables 5-7 from Mann et al. 2015) and  \nFigure 12. Absolute G-band magnitude plotted against estimated mass for 180 well studied late-K and M dwarfs from Mann et al. (2015). The solid blue line shows our fitted relation from Equation 18. See text for details.  \n<!-- image -->  \nrefined their mass estimates using more precise Gaia DR3 parallaxes and the Mann et al. (2019) mass-vs.MKs relation from above. They derived a relation between this estimated mass and the absolute G-band magnitude. However, the coefficients in Chontos et al. (2021) are published with insufficient accuracy to re-create the relation show in their figure 7, so we have re-derived them here. Our methodology is identical to theirs except that we exclude the sdM3 object L 750-42 (Gizis 1997) and do not incorporate a dependence on metallicity because the metallicity has not been measured for most of the M dwarfs within 20 pc. Using a functional form of  \nMass = 4 ∑ i =0 ci ( MG -10 . 5) i , (18)  \nwhere MG is the G-band absolute magnitude in magnitudes and Mass is in units of M ⊙ , we find best-fit coefficients of c 0 = 0 . 30548, c 1 = -0 . 10588, c 2 = 0 . 011471, c 3 = 0 . 0021352, and c 4 = -0 . 00041023. Our fit is illustrated in Figure 12. The relation is valid from 7 . 5 ≤ MG ≤ 15 . 0 (spectral types from ∼ K7 to ∼ M8). For uncertainty propagation, we adopt the Chontos et al. (2021) practice of a 2.2% uncertainty added in quadrature to the ∼ 3% uncertainty inherent to the Mann et al. (2019) relation. This massMG relation is particularly useful for estimating masses of individual components of close double systems that are currently resolved only by Gaia. In Table 4 we provide columns labeled "EstMassMG" and "EstMassMGErr" listing mass estimates for all objects for which these MG -based estimates can be computed.  \nThe second alternative estimation method is StarHorse (Anders et al. 2022), which uses Gaia EDR3 data crossmatched to photometry from Pan-STARRS1, SkyMapper, 2MASS, and AllWISE to estimate stellar parameters from stellar isochrones (from PARSEC 1.2S; Marigo et al. 2017) providing the closest match. When Anders et al. (2022)  \nmass estimates are available, these are listed in columns "EstMassSH" and "EstMassSHErr" of Table 4. These published mass uncertainties can be anomalously low compared to the other estimates discussed in this section because they pertain only to the internal model errors and do not include the systematic component coming from a model-to-truth comparison.  \nFigure 13 shows the four estimation techniques compared to each other. The top three panels show the intercomparisons between the TESS Input Catalog estimates, the MKs estimates, and the MG estimates. As these are all based on the same underlying mass vs. MKs relation of Mann et al. (2019), the correspondence is generally excellent. (In fact, the correspondence between the TESS Input Catalog estimates and estimates from our MKs technique are nearly perfect, differing only in the Gaia data release from which the parallax values were obtained.) The only deviation is for masses greater than ∼ 0 . 65 M ⊙ in the comparison between the TESS Input Catalog values and those derived from MKs , where the difference can be as large as 13%.  \nThe bottom three panels of Figure 13 show small systematics between the three estimation techniques above and StarHorse. As stated above, masses from StarHorse are based on theoretical models, so such systematics might be expected between theory and observation. At masses of ∼ 0 . 3 M ⊙ , StarHorse tends to overpredict (by ∼ 10%) the mass relative to the other techniques, and at smaller masses may significantly underpredict (by ∼ 35%). At masses near 0 . 8 M ⊙ , a small underprediction (by < 5%) relative to the TESS Input Catalog becomes an overprediction (by ∼ 5%) relative to masses from the MKs relation. At masses closer to 1 . 0 M ⊙ , StarHorse leads to underpredictions (by ∼ 10%) relative to estimates from the TESS Input Catalog.  \nGiven that systematic offsets of up to 15% are seen even between the sets with empirical underpinnings, we are reluctant to apply corrections to offsets smaller than this value. The only exception to this is the ∼ 35% offset seen for StarHorse estimates below StarHorse values of ∼ 0 . 275 M ⊙ . In this case, rather than applying an offset, we will simply not use any StarHorse estimates below 0 . 275 M ⊙ .', '7. FURTHER ANALYSIS': 'For each individual object ("#CompsOnThisRow" = 1) in Table 4, we have adopted a mass and its uncertainty. These are listed in columns "AdoptedInitialMass" and "AdoptedInitialMassErr" along with an additional column, "AdoptedInitialMassNote", indicating the origin of the data from elsewhere in the table. These are labeled with the term "Initial" as a reminder that for white dwarfs, we need their initial masses on the main sequence; for all other objects, their current masses are assumed identical to their initial masses. The codes for "AdoptedInitialMassNote" are as follows, listed in their order of selection:  \n- · wd IFMR, wd low, wd ultra-low , or wd conjecture : The initial mass and its uncertainty have been computed via  \nthe initial-to-final mass relation or other means (see Table 10), if this object is a white dwarf.  \n- · measured : Directly measured mass values from "Mass" and "MassErr" are used. The methodology used and its reference are listed in columns "MassMethod" and "MassRef". (For L, T, and Y dwarfs, directly measured masses are not retained because these are estimated in bulk through statistical means; see the Teff bullet, below.)\n- · M\\_Ks : The mass and its uncertainty from the Mann et al. (2019) MKs relation ("EstMassMKs" and "EstMassMKsErr") are used.\n- · TIC : The mass and its uncertainty from the TESS Input Catalog (Stassun et al. 2019; "EstMassTIC" and "EstMassTICErr") are used.\n- · M\\_G : The mass and its uncertainty from the MG relation of Equation 18 ("EstMassMG" and "EstMassMGErr") are used.\n- · SH : The mass and its uncertainty from StarHorse (Anders et al. 2022; "EstMassSH" and "EstMassSHErr") are used, unless that estimate falls below 0 . 275 M ⊙ (see section 6.3).\n- · literature : The mass and its uncertainty are taken from columns "EstMassLit" and "EstMassLitErr", the mass estimation method and reference for which are listed in "EstMassLitMethod" and "EstMassLitRef". (Literature values can supersede other values above if the object is listed as a giant or subgiant in Table 11.)\n- · see GeneralNotes : For objects with this code, the mass and its uncertainty were computed by us, as detailed in the "GeneralNotes" column of the table.\n- · Teff : For objects of type L, T, or Y, individual masses are not computed. These are handled statistically via the distribution of T eff values and their uncertainties ("Teff" and "Teff\\_unc"), as described in detail below.  \nFor cases in which literature values did not list a mass uncertainty, a value of 10% is arbitrarily assumed. The quoted StarHorse uncertainty is also replaced with a 10% uncertainty, based on the under- and over-predictions noted when comparing StarHorse values to other estimates (see discussion at end of Section 6.3), unless the quoted StarHorse internal uncertainty is already larger, in which case we retain the published value.  \nFor cases where only a miscellaneous magnitude or delta magnitude of a companion were available, it is instructive to estimate a spectral type for the object in order to estimate its mass. Figure 14 shows a comparison between masses and measured spectral types for those Table 4 objects having mass estimates (or direct measures) from one of the other methods. The piecewise fit shown in the figure is the one we  \nFigure 13. Intercomparisons of results from our four mass estimation techniques. The line of one-to-one correspondence is shown by the blue dashes. See text for details.  \n<!-- image -->  \nTable 16. Piecewise Fit to Mass vs. Dwarf Spectral Type Relation  \nNOTE-Each row in this table represents an inflection point in the red, piecewise fit of Figure 14.  \nuse to translate a dwarf spectral type estimate into a mass estimate. Other per-object details can be found in the General Notes column of Table 4.', '7.1. Analysis of brown dwarfs': 'We use the methodology adopted by Kirkpatrick et al. (2019, 2021a) and Raghu et al. (submitted) to determine the mass function for L, T, and Y dwarfs, most of which are brown dwarfs lacking any color (or spectral type or absolute magnitude) to mass correlation. Specifically, the mass function for these objects is determined by comparing the distribution of present-day temperatures to predicted temperature distributions. Predictions are drawn from a grid of models with varying mass functions, birthrates, and low-mass cutoffs. For each point in the grid, we build a predicted mass/age distribution that is then passed through a set of evolutionary models to predict the current-day T eff distribution. Using this grid of predictions allows us to find the combination of mass function, birthrate, and cutoff mass that best fits the observed temperature distribution.  \nFor the empirical distribution, we estimate the T eff value for each L, T, or Y dwarf (see Table 12 in this paper and table 11 of Kirkpatrick et al. 2021a) and then calculate space densities as a function of T eff. To compute space densities, we need to determine the distances at which our brown dwarf subsamples are truly complete, as the coldest Y dwarfs are so intrinsically dim that we are unable to push their completeness to the 20-pc limit targeted in this paper. As described in Kirkpatrick et al. (2021a), we determine completeness via the V / Vmax test (Schmidt 1968) using 150K bins and computing ⟨ V / Vmax ⟩ at half-parsec steps within each bin. The  \nFigure 14. Mass as a function of spectral type for 20-pc objects with measured (black points) or estimated (grey points) masses and optical spectral types in Table 4. The adopted initial mass (see text) is used for each object. For objects with estimated (not measured) masses, a random value between -0.25 and +0.25 has been added to the spectral type to better visualize otherwise overlapping data points. Our piecewise fit to the relation for dwarf stars is shown by the solid red line and is quantified in Table 16. For comparison, we show the average mass per spectral type as tabulated in the 2022.04.16 version of https://www.pas.rochester.edu/~emamajek/ EEM\\_dwarf\\_UBVIJHK\\_colors\\_Teff.txt (Pecaut & Mamajek 2013; magenta dashed line).  \n<!-- image -->  \ncomputation starts with the first half-parsec step falling just larger than the distance of the bin\'s nearest object and advances in distance out to d = 20 pc. These results are shown in Figure 15.  \nA comparison of this figure to figure 23 of Kirkpatrick et al. (2021a) shows that, despite the many new discoveries (and many fewer retractions) noted in Table 12, each 150K bin has the same completeness limit as before. As one example, consider the bin with the largest change, 600-750K. In both Kirkpatrick et al. (2021a) and here, this bin is complete out to 20 pc, but the number of objects has nonetheless increased from eighty-three in Kirkpatrick et al. (2021a) to ninety-eight in this paper; see also Table 17. (The V / Vmax test is only as robust as the Poisson statistics allow, which is why both  \nsets of numbers were deemed to be complete.) As another example, the number of objects interior to the completeness limit of 15.0 pc in the 450-600K bin has increased from fiftythree to fifty-six.  \nAs noted in Kirkpatrick et al. (2021a), the V / Vmax test does not check for inhomogeneities in surface area, the most likely cause of which would be confusion along the Galactic plane that hinders our ability to find nearby brown dwarfs. Do the increased densities now reported in this paper indicate that these corrections can be reduced or dropped altogether?  \nFigure 16 shows the positions in Galactic coordinates of all 583 L, T and Y dwarfs in the 20-pc census. As was done in Kirkpatrick et al. (2021a), we divide the sky into two zones: a zone along the Galactic plane ( | glat | < 14 . · 48) and another ( | glat | ≥ 14 . · 48) well outside of the plane. This value of | glat | was chosen so that the non-plane zone contains exactly three times the area of the plane zone. If there is no incompleteness along the Galactic plane, then the ratio of non-plane to plane objects should be three. For the volumecomplete portions of our 20-pc census, we find that this ratio is 138 / 44 = 3 . 1 for L dwarfs, 257 / 65 = 4 . 0 for T dwarfs, and 31 / 4 = 7 . 8 for Y dwarfs, suggesting that the Galactic plane does not introduce any significant incompleteness ( < 1%) for L dwarfs but does still impede the discovery of fainter T and Y dwarfs. In contrast, Kirkpatrick et al. (2021a) derived ratios of 137 / 34 = 4 . 0, 234 / 34 = 6 . 9, and 24 / 4 = 6 . 0 for the L, T and Y dwarf samples, respectively.  \nIncompleteness along the Galactic plane has improved in the current 20-pc census for the L and T dwarfs. For the Y dwarfs, the view is complicated by smaller number statistics. Taking the non-plane numbers of Y dwarfs as truth, then the number of plane Y dwarfs in the current sample should be 31 / 3 ± ( √ 31) / 3 = 10 . 3 ± 1 . 9, which is 3 . 3 σ different from the value of 4 actually found. The same computation for the Kirkpatrick et al. (2021a) numbers gives a number of plane Y dwarfs that was only 2 . 5 σ different. Hence, the underdensity of Y dwarfs in the plane is now significantly worse, due to the fact that all new discoveries of Y dwarfs within the volume have been found outside of the plane zone.  \nL dwarfs no longer show an underdensity in the plane, so no correction is needed for our derived L dwarf space densities. T dwarf space densities should, however, be multiplied by 1.06 to account for the observed incompleteness. The Y dwarf incompleteness is harder to assess given the small number of Y dwarfs in the plane, but the raw numbers suggest a conservative correction factor of 1.15, slightly larger than the 1.13 factor adopted by Kirkpatrick et al. (2021a). These factors are listed in Table 17.  \nThe final step in measuring the space densities of L, T, and Y dwarfs is assessing their measurement uncertainties. For this we adopt the same methodology used in Kirkpatrick et al. (2021a). To summarize, our confidence in assigning an object to a T eff bin is directly related to the measurement uncertainty on T eff, which is often comparable to the bin size itself. To estimate our confidence in the numbers of objects in each bin, we have run simulations with 10,000 Monte  \nFigure 15. The average V / V max value in 0.5-pc intervals across fourteen 150-K bins encompassing L, T, and Y dwarfs. Blue dots show the empirical sample, and red labels denote the number of objects at each 0.5-pc computation. The black dashed line shows the ⟨ V / Vmax ⟩ = 0 . 5 level indicative of a complete sample. The grey error bars show the approximate 1 σ range that a sample of the size shown in red would exhibit, given random statistics. The brown error bars, offset by +0.05 pc from the grey error bars for clarity, show the 1 σ variation obtained by simulations using 10,000 Monte Carlo realizations having the number of objects and completeness limit listed in Table 17. See section 8.2 of Kirkpatrick et al. (2021a) for more details.  \n<!-- image -->  \nFigure 16. Plots of the 20-pc L, T, and Y dwarf census in Galactic coordinates. The four panels display (a) the sample in its entirety (black), (b) only the L dwarfs (blue), (c) only the T dwarfs (green), and (d) only the Y dwarfs (red). New additions to the sample since Kirkpatrick et al. (2021a) are plotted with grey haloes in panels (b) through (d).  \n<!-- image -->  \nTable 17. Space Densities for Early-L through Early-Y Dwarfs  \nb This value is computed via the equations  \ndens = ( raw )( corr ) / ( 4 3 π dmax 3 )  \nand  \nFigure 17. Our measured L, T, and Y dwarf space densities from Table 17 (black dots) as a function of effective temperature overplotted on different simulations from Raghu et al. (submitted). In all panels, simulations assuming a constant birthrate are shown, along with the results for three different low-mass cutoffs: 10 MJup (light blue), 5 MJup (green), and 1 MJup (red). Panels in the left column use the Saumon & Marley (2008) evolutionary models, and panels in right column use Baraffe et al. (2003). The top row shows simulations with a power law of α = 0 . 8, the middle row shows α = 0 . 6, and the bottom row shows α = 0 . 4. The best overall fits are those shown in the left panel in the middle row, using α = 0 . 6 and the Saumon & Marley (2008) models.  \n<!-- image -->  \noutside the completeness limit of the colder bin. This last loss is one-sided, however, as any colder objects scattering into the warmer bin would be necessarily retained. Hence, we compute our adopted space densities using the raw number counts, but including the uncertainties derived from our simulations, as shown in the footnote of Table 17. These densities are graphically illustrated in Figure 17.  \nThe measured space densities can now be compared to the simulated T eff distributions (Raghu et al., submitted) to infer the form of the mass function at this low-mass end. Following on the results of Kirkpatrick et al. (2021a), which showed the best match to be a power law, dN / dM ∝ M -α , with α ≈ 0 . 6, Raghu et al. (submitted) assume power-law functional forms with α values between 0.3 and 0.8 and, like  \nσ dens = √ ( σ raw 2 + σ ad j 2 ) ( corr ) / ( 4 3 π dmax 3 )  \nwhere σ raw = √ raw .  \nCarlo realizations wherein we take the uncertainty in T eff and multiply it by a random value generated from a normal distribution having a mean of 0 and a standard deviation of 1. For each simulation, this uncertainty is added onto the measured value and the object (re-)assigned to the appropriate T eff bin. The computed means and standard deviations across all 10,000 realizations are given in column 5 of Table 17. We use only these computed standard deviations in our adopted space densities, but not the adjusted means. As further explained in Kirkpatrick et al. (2021a), the reason for this is that the number of objects is not preserved across the Monte Carlo simulations because some objects scatter into the hotter, incomplete bin at 2100-2250K and are lost, while objects at the other temperature extreme may be lost because they fall  \nKirkpatrick et al. (2021a), choose low-mass cutoffs of ∼ 1, 5, and 10 MJup . Unlike Kirkpatrick et al. (2021a), however, they vary the birthrate to include not only a constant birthrate over the lifetime of the Milky Way, but also consider two other birthrates - called inside-out and late-burst - from Johnson et al. (2021) that are constrained by new results from Gaia. The inside-out birthrate represents a declining birthrate over the 10 Gyr lifetime of the Galactic disk, and the late-burst birthrate is identical to the inside-out form, except with an abrupt increase (by a factor of ∼ 3) in star formation ∼ 3-5 Gyr ago.  \nEvolutionary models are used to infer the current T eff value of each simulated object (using its mass and age). Raghu et al. (submitted) expand the model set used in Kirkpatrick et al. (2021a) by including the newer Marley et al. (2021) predictions and show (again) that the only evolutionary models able to fit the bump in the L/T transition in the T eff distribution are those of Saumon & Marley (2008).  \nIt has been shown in Kirkpatrick et al. (2021a) and Raghu et al. (submitted) that the low-mass cutoff has little effect on the shape of the mass function at T eff values above 450K, where our fitting is taking place. Therefore, we consider each α + birthrate pair and compute the median of the least squared values for the simulations across all three cutoff masses. The minimum is achieved for α = 0 . 6 and a constant birthrate, identical to the findings in Kirkpatrick et al. (2021a). The second best fit is achieved for α = 0 . 5 and a constant birthrate. The third best fit is a tie among the α = 0 . 7 + constant, the α = 0 . 4 + late-burst, and α = 0 . 5 + late-burst models. Use of either the late-burst or inside-out birthrates results in a slightly reduced α because those birthrates create a small overabundance, relative to the constant birthrate models, of older brown dwarfs that have already cooled to cooler temperatures.  \nIn Figure 17, we show the fits for three values of α (0.4, 0.6, and 0.8) all paired with a constant birthrate. The panels in the left column of the figure show that the α = 0 . 6 model with a constant birthrate and using the Saumon & Marley (2008) evolutionary models is an excellent representation of the empirical data. Can any new conclusions be gleaned regarding the low-mass cutoff? As figure 4 of Raghu et al. (submitted) illustrates, the Saumon & Marley (2008) models are incomplete below masses of ∼ 0.015 M ⊙ ( ∼ 16 MJup ), so they are a poor choice for determining what the low-mass cutoff might be. Instead, we revert back to the Baraffe et al. (2003) models, that are complete down to ∼ 5 MJup . As the rightmost panels in Figure 17 illustrate, our ability to distinguish between low-mass cutoffs depends on measuring accurate space densities below 450K. Using the 20-pc census to say confidently that star formation\'s terminus is below 10 MJup or even 5 MJup depends on surveying the sky more deeply at the wavelengths of these objects\' peak emission and obtaining the necessary astrometry to measure accurate distances. As the simulations using the Baraffe et al. (2003) models show, measuring an accurate space density for the 300-450K bin will allow us to distinguish between the cutoff masses, and even a few more objects discovered in the  \nFigure 18. Plots of the implied space densities of brown dwarfs (brown) in 0.001 M ⊙ bins compared to the measured space densities of other possible low-mass star formation products from Table 15 (blue). The brown dwarf space densities are divided into three mass zones M > 10 MJup (solid brown), 5 MJup < M < 10 MJup (dashed brown), and 1 MJup < M < 5 MJup (dotted brown). Note that the densities of the possible pseudo-exoplanets do not affect our measurement of the brown dwarf space densities, as their numbers only become appreciable at masses well below 5 MJup .  \n<!-- image -->  \n150-300K bin, which currently has only the 250K Y dwarf WISE J085510.83 -071442.5 in it, will provide even tighter constraints.  \nHave some of these ultra-low mass products of star formation already been identified, and are they masquerading in the literature as exoplanet discoveries to higher mass objects? We use our analysis in Table 15 to see first if the omission of these objects has biased our derivation of the brown dwarf mass function above. With the exception of the two objects (the companions to UCAC4 211-005570 and L 119213) lacking mass estimates, we take all objects labeled as "consider" in column 4 of Table 15 and estimated their contribution to the overall mass function. For objects with M sin( i ) measurements only, we pull a random number from a distribution of values uniformly distributed between 0 and 1 and multiply that number by 90 degrees to assign each an inclination, which we then use to assign an actual mass value. For all masses, whether or not they are true masses or adjusted M sin( i ) measurements, we then pull a random number from a normal distribution with a mean of 0 and standard deviation of 1 and multiply that number by the uncertainty, which we then add back to the mass value. We perform this methodology over 10,000 Monte Carlo iterations and find the mean and standard deviation of the resulting space density, binned over 0.001M ⊙ mass intervals, as illustrated in Figure 18.  \nThis figure shows that our derived space density of brown dwarfs, (which we find to be ξ ( M ) = dN / dM = 0 . 0469 × M -0 . 6 in units of # pc -3 [1 M ⊙ ] -1 , with M in units of M ⊙ ; see  \nSection 8), overwhelms the space density above 5 MJup where our fitting took place. So, the omission of these objects has no impact on our derivation. However, the second question is whether any of these objects could be products of star formation itself rather than the secondary by-products of a protoplanetary disk. That question cannot be answered from this diagram, but it is a statistical certainty that at least a few of the objects on the high-mass tail of this distribution are star formation products. One striking result from Figure 18, however, is the high space density of objects in the lowest mass bin, given that the census of such low-mass objects, whether resulting from star formation or proto-planetary disks, is still woefully incomplete. There is clearly no shortage of ultralow mass objects in the Milky Way.', '7.2. Combined stellar and brown dwarf space densities': "With the brown dwarf portion of the mass function now fitted, we can combine the stellar and brown dwarf portions to determine the shape of the overall mass function.  \nFirst we take the number counts across the stellar regime and perform a similar Monte Carlo analysis as was done on the brown dwarfs. Specifically, for each object we pull a random number from a normal distribution with a mean of 0 and a standard deviation of 1. We then multiply the mass measurement uncertainty by the random number and add that back to the mass value to get a true mass. We do this for each of the stars in our sample, and repeat the process 10,000 times to simulate 10,000 possible histograms. We then compute the mean value in each histogram bin along with its standard deviation. Because the 20-pc volume around the Sun is just one of many such volumes that can be taken as a sample of the Milky Way, we add the Poisson uncertainty and the standard deviation from above in quadrature to provide the final uncertainty per bin. (This parallels the brown dwarf space density analysis of Table 17.)  \nWe can now append the substellar contribution onto this stellar distribution. To do this, we look at the predictions from the best fit Raghu et al. (submitted) model to the brown dwarf T eff distribution from above, which is the α =0 . 6 power law with a constant birthrate function and passed through the Saumon & Marley (2008) evolutionary models. We also choose a 0.005 M ⊙ ( ∼ 5 MJup ) cutoff to parallel the more detailed cutoff analysis from Kirkpatrick et al. (2021a). This simulation gives the predicted mass distributions shown in Figure 19. Each histogram is scaled so that the total number of objects in each histogram matches the raw numbers of objects per bin listed in Table 17. As one example, the 27 objects in the 1950-2100K bin are predicted to fall almost exclusively in the 0.075-0.080 M ⊙ bin, and these predictions suggest that our 27 objects be apportioned as 20.5 objects in the 0.075-0.080 M ⊙ bin, 2.0 objects in the 0.0700.075 M ⊙ bin, 1.0 object in the 0.065-0.070 M ⊙ bin, and fractional numbers of objects in bins of lower mass. As another example, the 63 objects in the 750-900K bin are spread over a wide range of masses from 0.005-0.060 M ⊙ and are apportioned as 2.6 objects in the 0.055-0.060 M ⊙ bin, 10.8 objects in the 0.050-0.055 M ⊙ bin, 12.3 objects in the 0.045-  \n0.050 M ⊙ bin, 10.7 objects in the 0.040-0.045 M ⊙ bin, 8.3 objects in the 0.035-0.040 M ⊙ bin, etc.  \nWe take the apportionment across all thirteen temperature bins and tally the results in each of the 0.005 M ⊙ -wide mass bins, after also applying the factor ( corr in Table 17) to correct for losses of objects along the Galactic plane and extrapolating the numbers to the full 20-pc volume if that temperature bin was not complete to 20 pc. For example, the raw number counts in the 450-600K bin shown in both Table 17 and Figure 19 were multiplied by the 1.06 correction factor then multiplied by (20 / 15) 3 to extrapolate to the full volume. In Table 17, we find that our lowest temperature bin with a space density measurement, 300-450K, is considered to be incomplete, and the mass distribution for that bin in Figure 19 suggests that that bin's objects fall exclusively below 0.025 M ⊙ . Therefore we consider any space density measurements below this mass value to be lower limits only.  \nWe now add these brown dwarf masses to the results of our Monte Carlo analysis of stellar masses above to produce a mass function across the entire mass range. This initial mass function is illustrated in Figure 20. Panel (a) shows the mass function across the full mass range from 0 to 8 M ⊙ , binned in 0.1 M ⊙ increments. The mass function rises with decreasing mass, and it continues to rise beyond our 0.025 M ⊙ ( ∼ 26 MJup ) completeness limit. Subsequent panels show details. Panel (b) shows the high-mass end of the initial mass function from 1.5 to 8.0 M ⊙ , again with 0.1 M ⊙ binning. The statistics above 3 M ⊙ are poor but nonetheless show a steady increase from there down to 1.5 M ⊙ . Panel (c) shows the mid-mass range (0 . 4 < M < 1 . 5 M ⊙ ), now binned into smaller 0.02 M ⊙ increments because the statistics here are richer. Panel (d) zooms in on the smallest mass portion, below 0.4 M ⊙ , and chooses yet a smaller mass binning of 0.005 M ⊙ . With the exception of a few small features (discussed below), the mass function is seen to rise monotonically from 1.5 to 0.025 M ⊙ . Mostly within the measurement errors (see more discussion below), the initial mass function is seen to continue rising well below our 0.025 M ⊙ completeness limit and at least down to 0.015 M ⊙ .  \nThe numbers on which Figure 20 is based are given in Table 18. For ease of reference, both the number of stars and the space density is given for each mass bin. Three mass binnings are tabulated, roughly paralleling what is shown in Figure 20: 0.1 M ⊙ binning across the entire 0.0-8.0 M ⊙ range (80 bins), 0.02 M ⊙ binning across the range 0.0-1.6 M ⊙ (80 bins), and 0.005 M ⊙ binning across the range 0.0-0.4 M ⊙ (80 bins). Mass bins with incomplete statistics are labeled as lower limits in the final column of the table.  \nThere are a few features in Figure 20(c) and (d) that warrant special attention. The first is the bump in the object counts near 0.55 M ⊙ in panel (c). This falls near the point at which our mass estimation switches from that of the TIC relations of Stassun et al. (2019) at higher masses to that of the MKs relation of Mann et al. (2019) at lower masses. Currently, we switch between these two relations at MKs = 5 . 0 mag, corresponding to a mass of ∼ 0 . 6 M ⊙ . As a test, if we  \nSun  \nFigure 19. The predicted distributions of brown dwarf masses in each of our 150K effective temperature bins based on the best-fit Raghu et al. (submitted) simulation to our measured L, T, and Y dwarf space densities (see text for details). Each histogram is scaled to match the total number of objects listed for that T eff bin in Table 17.  \n<!-- image -->  \nTable 18. Number of Objects and Space Densities per Mass Bin for the 20-pc Census  \nNOTE-(This table is available in its entirety in machine-readable form.)  \n- a The [ M bin] -1 portion of the units should be replaced with the bin size for that row. For example, for the first row of the table, the units will be pc -3 [0 . 10 M ⊙ ] -1 because that bin is 0 . 10 M ⊙ wide.  \nchange the switchover point to be at MKs = 4 . 0 mag ( M ≈ 0 . 7 M ⊙ ) instead, we find that this bump in the space densities moves to higher masses, with a deficit around 0 . 8 M ⊙ , as shown in Figure 21. We also note that the uncertainties in the masses resulting from the Mann et al. (2019) relation are three to four times smaller than those derived from the Stassun et al. (2019) relation. As another test, we can artificially inflate the mass uncertainties on the Mann et al. (2019)derived masses while keeping the current switchover point at MKs = 5 . 0 mag. That result is also shown in Figure 21. In this case, the bump is greatly diminished in the number counts, but an inflection point is still seen near 0.6 M ⊙ . Given that this feature in the number counts moves in response to the mass estimation used, we believe it is an artificial effect. Furthermore, given that the Stassun et al. (2019) mass estimation relies on dynamically measured individual masses whereas the Mann et al. (2019) relation uses Bayesian statistics to ferret out individual masses from binaries in which only the total system mass is measured, this likely indicates a small systematic offset that slightly deflates the Mann et al. (2019)-derived masses relative to truth. In fact, an effect in this direction and representing a systematic offset of ∼ 2% is seen when comparing results of the MKs relation to individually derived masses (figure 15 of Mann et al. 2019). Obtaining more directly measured individual masses in this regime, corresponding to late-K and early-M dwarfs, would help to put this issue to rest.  \nOther features are seen in Figure 20(d). There is a small drop in the number counts near 0.13 M ⊙ followed by a sudden rise near 0.11 M ⊙ . This feature is likely artificial, as the bump at 0.11 M ⊙ is due primarily to components in multiple systems about which little information is known, and these were arbitrarily assigned masses of 0.11 M ⊙ based on an an-  \n0.00  \nFigure 20. The 20-pc initial mass function across all stellar and substellar masses. Our measured values and their uncertainties are shown in black. The raw number counts for stars of type M9.5 and earlier are shown by the blue histogram. (a) The full mass range, 0.0-8.0 M ⊙ , with 0.1 M ⊙ binning; (b) A zoom-in of the high-mass end, from 1.5 to 8.0 M ⊙ , with the same binning; (c) A zoom-in of the mid-range, from 0.4 to 1.5 M ⊙ , with 0.02 M ⊙ binning. (d) A zoom-in of the low-mass portion, 0.0-0.4 M ⊙ , with 0.005 M ⊙ binning. Our fit to the mass function is shown by the orange line.  \n<!-- image -->  \nticipated spectral type of M5.5. This type lies at a very sharp inflection point (see Table 16) in our mass vs. spectral type relation (Figure 14).  \nThe other feature is the rapidly changing number count between masses of 0.06 and 0.08 M ⊙ , a mass range that straddles the stellar/substellar break. Some of the early- to mid-L dwarfs that we have included in our brown dwarf mass function analysis are likely very low-mass stars and not brown dwarfs themselves. As a consequence, these are assigned masses that are a bit too high (the extraneous high point in the 0.075-0.080 M ⊙ bin), which likely leads to concomitant deficits in the next higher mass bins. In fact, Table 4 lists three early-L dwarfs in the 20-pc census that have dynamical mass measurements, and one of these (LP 388-55 B; Dupuy  \n& Liu 2017) has a mass just above the 0.075-0.080 M ⊙ mass bin (0.083 ± 0.03 M ⊙ ).", '8. DISCUSSION': 'How do our 20-pc results compare to other attempts in the literature to measure the initial mass function? Pioneering work by Salpeter (1955) found that a power law form ξ ( M ) = dN / dM ∝ M -α with α =2 . 35 best fit the initial mass function over the range 0 . 3 ≤ M ≤ 10 M ⊙ . Subsequent work by Miller &Scalo (1979) found α =1 . 4 for 0 . 1 ≤ M ≤ 1 M ⊙ and α =2 . 5 for 1 ≤ M ≤ 10 M ⊙ . Scalo (1986) determined α = 2 . 7 for 2 ≤ M ≤ 10 M ⊙ , and Reid et al. (2002) found α = 1 . 3 for 0 . 1 ≤ M ≤ 0 . 7 M ⊙ .  \nAs more accurate measurements of the initial mass function became possible, researchers realized the importance of  \n<!-- image -->  \n<!-- image -->  \nFigure 21. Tests of the bump in the initial mass function seen near 0.55 M ⊙ in Figure 20. (a) A zoom-in showing the bump in Figure 20(c). (b) The number counts over the same mass range but where we have moved the switchover in the mass estimate from MKs = 5 . 0 mag to MKs = 4 . 0 mag. (c) The number counts over the same mass range but where we have kept the MKs = 5 . 0 mag switchover point and inflated the uncertainties on the MKs -derived masses by a factor of four. For other details, see the caption to Figure 20.  \n<!-- image -->  \ndistinguishing whether the masses used for the computation were that of the stellar system or of its individual components. For example, in an analysis of data from the Sloan Digital Sky Survey, Bochanski et al. (2010) found α = 2 . 38 for 0 . 32 < M < 0 . 80 M ⊙ and α = 0 . 35 for 0 . 10 < M < 0 . 32 M ⊙ for the mass function of systems but α = 2 . 66 and α = 0 . 98 for the mass function of single stars over the same two mass regimes, respectively. Earlier, Kroupa (2001) had found that single-star initial mass functions resulting from the analysis of young star clusters generally gave values of α that were higher by ∼ 0.5 (for 0 . 1 < M < 1 . 0 M ⊙ ) than the field initial mass function, for which systems were not resolved. Reid (2005) cautions that unresolved multiplicity complicates interpretation of the initial mass function; the initial mass function of systems is more directly tied to the fragmentation of the original molecular cloud, but the initial mass function of individual objects give the mass distribution of the (sub)stellar bodies formed.  \nOur work on the 20-pc census has concentrated on analysis of the individual products of star formation, as we are curious to know how frequently this process produces, for example, very low-mass brown dwarfs compared to higher-mass objects. We will therefore restrict subsequent analysis here to the single-object initial mass function and leave analysis of the 20-pc census regarding multiplicity and the system mass function to those investigating how the formation of systems relates back to cloud fragmentation.', '8.1. Comparison of Initial Mass Functions': "Here we compare two very well established forms of the initial mass function and compare their predictions to our results based on the 20-pc census.  \nChabrier (2001, 2003a,b) developed several functional forms for the initial mass function. The latest one relating to single objects is given by Chabrier (2003b) and is comprised of a power law at high masses and a log-normal form  \nat lower masses:  \nξ ( M ) = C 1 ln10 M -α , for M > 1 . 0 M ⊙ = C 2 M ln10 e -(log M -log Mc ) 2 2 σ 2 , for M ≤ 1 . 0 M ⊙ (19)  \nwhere C 1 = 0 . 0443, α = 2 . 3, C 2 = 0 . 158, Mc = 0 . 079 M ⊙ , and σ = 0 . 69 M ⊙ . The value of ξ ( M ) is in units of # of objects per pc 3 per M ⊙ . See equation 2 and table 1 of Chabrier (2003b) for the derivation shown above. We note that the values of C 1 and C 2 are set by Chabrier (2003b) to match an empirical space density measurement (at 1 M ⊙ from Scalo 1986) of the initial mass function in the Milky Way's disk population.  \nLikewise, Kroupa et al. (2013) found that a tripartite power-law form best describes the single-object initial mass function:  \nξ ( M ) = C ( 0 . 5 0 . 07 ) -α 2 ( M 0 . 5 ) -α 1 , for 0 . 5 < M < 150 M ⊙ = C ( M 0 . 07 ) -α 2 , for 0 . 07 < M < 0 . 5 M ⊙ = C 3 ( M 0 . 07 ) -α 3 , for 0 . 01 < M < 0 . 15 M ⊙ (20)  \nwhere α 1 = 2 . 3, α 2 = 1 . 3, and α 3 = 0 . 3. As above, the value of ξ ( M ) is in units of # of objects per pc 3 per M ⊙ . Note that there are two components of this mass function that contribute to the mass range 0 . 15 < M < 0 . 07 M ⊙ . The value of C is not specified by Kroupa et al. (2013). However, figure 4-24 of Kroupa et al. (2013) provides a comparison of this initial mass function with the Chabrier (2003b) version in Equation 19, showing that they are identical at ∼ 0.85 M ⊙ , resulting in a value of C ≈ 1 . 15.  \nAs shown in Figure 22, neither of these parameterizations adequately describes the 20-pc initial mass function derived here. The main reason for this is that prior determinations were done pre-WISE, pre-Spitzer, and pre-Gaia before the L,  \nT, and Y dwarf complement of the mass function was fully characterized and before exquisite parallax determinations became available for almost all objects in the 20-pc volume. As such, these prior works relied on poorer statistics with fits done in logarithmic scaling on both the dN and dM axes. Our careful accounting of objects within the 20-pc volume now allows for a more precise determination of the single-object initial mass function.  \nWe thus provide a new multi-part power law parameterization that is bounded by the following caveats: (1) We assume α = 2 . 3 at the high mass end, as has been determined from earlier studies, and we do this because our 20-pc census has few stars with M > 2 M ⊙ to better constrain this. (2) We assume α = 1 . 3 at intermediate masses, as this has also been established by earlier studies. (3) We take α = 0 . 6 in the brown dwarf regime, as was determined in Section 7.1. We do not constrain the stitch points in mass between the power law pieces nor do we limit the number of power law pieces to only three. We perform these fits by eye, keeping in mind the caveats from Section 7.2 concerning the non-physical bumps and troughs in the number counts as a function of mass. Given the constraints above, we find that a three-piece power law does not adequately describe the number counts in the mid- to late-M dwarf regime (0 . 08 ≲ M ≲ 0 . 20 M ⊙ ), but that a four-piece power law can. This best fit is given below and illustrated by the orange line in Figure 20:  \nξ ( M ) = C 1( M ) -α 1 , for 0 . 55 < M < 8 . 0 M ⊙ = C 2( M ) -α 2 , for 0 . 22 < M < 0 . 55 M ⊙ = C 3( M ) -α 3 , for 0 . 05 < M < 0 . 22 M ⊙ = C 4( M ) -α 4 , for 0 . 01 < M < 0 . 05 M ⊙ (21)  \nwhere C 1 = 0 . 0150, α 1 = 2 . 3, C 2 = 0 . 0273, α 2 = 1 . 3, C 3 = 0 . 134, α 3 = 0 . 25, C 4 = 0 . 0469, α 4 = 0 . 6, As above, the value of ξ ( M ) is in units of # of objects per pc 3 per M ⊙ .  \nIf we integrate under our best fit from 0.075 to 8.0 M ⊙ , we find 3002 stars, which can be compared to the 3000 individual objects in the 20-pc census that have (measured or implied) types of M9.5 or earlier. The integration under our fit for 0.020 to 0.075 M ⊙ gives 789 brown dwarfs, compared to the 582 individual L, T, and Y dwarfs in Table 4. Most of the missing ∼ 200 brown dwarfs are ones with T eff values 450-600K and distances of 15-20 pc or ones with 300-450K temperatures and 11-20 pc distances, where our current accounting is known to be incomplete (see Table 17). These results show that the number of stars relative to brown dwarfs is 3002/789, or ∼ 4. However, we believe that the brown dwarf mass function extends to at least 0.010 M ⊙ , which would give a ratio of 3002/986 ( ∼ 3) if the α = 0 . 6 functional form continues to that mass. In the limiting case in which it continues to zero mass, we find a star-to-brown-dwarf ratio of 3002/1602, or ∼ 2. We note that as recently as a decade ago, this ratio in the Solar neighborhood was believed to be as high as 6:1 (Kirkpatrick et al. 2012).  \nThis decrease in the ratio of stars to brown dwarfs is not in tension with microlensing results, as analysis of OGLE data found that an even steeper power law in the brown dwarf  \nregime ( α = 0 . 8 for 0 . 01 < M < 0 . 08 M ⊙ ; Mróz et al. 2017) best fits the observed distribution of short-timescale (lowmass) events. Furthermore, the possibility that the mass function extends into a regime significantly below 0.010 M ⊙ is bolstered by recent JWST observations of the Orion Nebula Complex that show a significant population of objects, down to at least 0.001 M ⊙ , that are apparently direct products of star formation (McCaughrean & Pearson 2023; Pearson & McCaughrean 2023).  \nFinally, we note that our accounting of the mass of hydrogen-burning stars in the 20-pc census along with our best fit to the mass function of brown dwarfs allows us to calculate the average mass of an object in this sample. Integrating our mass function down to a mass of 0.020 M ⊙ , we find that value to be 0.41 M ⊙ . There are likely many undiscovered brown dwarfs in the solar neighborhood too faint to be currently detected, so this average value could be pushed lower. Assuming there is no low-mass cutoff of star formation and the brown dwarf mass function continues to zero mass with a power law slope of α = 0 . 6, we find that the average mass of objects in the 20-pc census would drop to 0.34 M ⊙ . This can be considered as the limiting case unless the functional form at the lowest masses has an α value considerably greater than 0.6.", '9. CONCLUSIONS': 'In this paper our aim was to study the initial mass distribution of star formation\'s by-products, from the highestmass progenitors of present-day white dwarfs to the lowestmass brown dwarfs. For this, we have produced a volumecomplete sample of stellar and substellar objects within 20 pc of the Sun. We have split multiple systems into their separate components and characterized each individual object to provide an accurate mass assignment.  \nOur main conclusions can be summarized as follows:  \n1) The initial mass function steadily increases as a function of descending mass. Its peak in (linear) mass space is not yet defined but is located below 0.020 M ⊙ ( ∼ 20 MJup ; Figure 20). We find that a quadripartite power-law ( ξ ( M ) = dN / dM ∝ M -α ) fits the observed space densities well (Equation 21). Going from high mass to low mass, we find exponents of α = 2.3, 1.3, 0.25, and 0.6, with stitch points between segments of 0.55, 0.22, and 0.05 M ⊙ , respectively. Although the rate of ascent of the mass function is slowly retarded as a function of descending mass through the stellar and high-mass brown dwarf regimes, its ascent increases again for the lower-mass brown dwarfs.  \n2) This initial mass function agrees well with previous determinations in the high-mass regime (by design) but differs markedly from other established formalisms in the M, L, T, and Y dwarf regimes (Section 8). Functional forms proposed by Chabrier (2003b) and Kroupa et al. (2013) overpredict the number of these lower-mass dwarfs relative to their more massive counterparts.  \n3) The 20-pc census currently consists of ∼ 3000 stars and ∼ 600 brown dwarfs (Table 4). At face value, this implies a stellar-to-substellar ratio of ∼ 5, but corrections for in-  \n<!-- image -->  \n<!-- image -->  \nFigure 22. A comparison of our 20-pc number counts (black points with error bars) and our fit of the initial mass function (solid orange line) to the functional forms of Kroupa et al. (2013) (dotted blue line) and Chabrier (2003b) (dashed green line). Each panel shows a zoom-in of a different mass segment: (a) 1 . 5 < M < 8 . 0 M ⊙ , (b) 0 . 4 < M < 1 . 5 M ⊙ , (c) 0 . 0 < M < 0 . 4 M ⊙ .  \n<!-- image -->  \nteness for brown dwarfs with temperatures from 300600K shows that the ratio is currently measured at ∼ 4. Incompletenesses at lower temperatures may yet bring this ratio as low as ∼ 3 (Section 8). The average mass of objects in the 20-pc census is currently measured as 0.41 M ⊙ but could drop as low as 0.34 M ⊙ if many colder brown dwarfs, yet to be discovered, actually exist (Section 8).  \n4) The 20-pc census of objects colder than 600K, corresponding to spectral type ∼ T8.5, is still incomplete beyond 15 pc, and the completeness volume shrinks to 11 pc at 450K, corresponding to spectral type ∼ Y0 (Table 17). Moreover, an additional source of incompleteness for objects as warm as 1350K is high backgrounds and confusion along the Galactic plane (Section 7.1).  \n5) There are direct indications that many unrecognized companions still exist to already identified members of the 20-pc census. Acceleration (aka proper motion anomaly) has been used to flag many such systems (Section 5.1.1), and large Gaia RUWE/LUWE values significantly higher than 1.0 flag many others (Section 5.2). Additional follow-up of these systems would help to better flesh out the 20-pc census itself while also providing much firmer statistics on multiplicity and the prevalence of hierarchical systems.  \n6) Our "complete" (see caveats #4 and #5, above) 20-pc census produced for this paper is available for additional uses (Table 4). As one example, this nearby sample is particularly useful as the hunting grounds for the closest habitable worlds to our own Solar System and is thus also available via the NASA Exoplanet Archive 44 .  \n7) Except for white dwarfs (Section 6.1.1) and brown dwarfs (Section 6.1.3), masses are used directly when they have been measured (Section 4). Most objects, though, lack actual mass determinations. For these we use a variety of  \nmass estimation methods and select the ones that provide the most reliable results, when comparison to truth is available (Section 6.3). Nonetheless, our resulting space density computations binned by mass show some spurious features that appear to be caused by shortcomings in the estimation method. These are most obvious in the early-M dwarf region and in the regime from late-M to early-L dwarfs (Section 7.2). Dedicated programs, such as those by Vrijmoet et al. (2022) and Dupuy et al. (2022) directly determining masses of objects in these zones are clearly needed.  \n8) Dueling definitions in the literature of "brown dwarf" and "exoplanet" could bias our results if objects tagged as exoplanets are omitted from the initial mass function. We account for this and find that most of the objects labeled exoplanets (via the 13 MJup -based definition) fall in a small-mass regime separate from the objects that have been more traditionally labelled as brown dwarfs (via the formation-based definition) and do not affect our conclusions regarding the mass function (Section 6.2.3). This having been stated, future studies of the initial mass function might wisely consider no such division, as planet formation can be thought of merely as a (delayed) secondary process resulting from star formation itself. Including planetary formation products as another branch of the initial mass function will, however, not be feasible until a statistically robust, volume-complete set of exoplanets can be reliably measured.  \n9) Gaia DR3 detections comprise only ∼ 75% of the volume-complete 20-pc census. Objects within 20 pc of the Sun can be missed by Gaia because they are too bright for Gaia observations, too faint for Gaia to detect, or are companions to Gaia-detected host stars that are (presently) inadequately characterized astrometrically (Section 2.1).  \n10) Citizen science continues to produce new discoveries within 20 pc (Section 2.1.2), even including a possible new Y dwarf with a "bonus" Spitzer parallax (Section A.3). Such discoveries are coming largely from WISE data sets, but these sets are being pushed to their sensitivity limits. Com-  \npleting the 20-pc census in the 300-600K temperature range will require a deeper survey at ∼ 5 um, the best prospect for which is the NASA mission NEO Surveyor (Kirkpatrick et al. 2019b; Mainzer et al. 2023). Results from that mission, along with searches for cooler companions to known 20-pc hosts with JWST, represent our best short-term prospects for determining the occurrence rate of objects such as WISE J085510.83 -071442.5 that reside below 300K.  \n11) The 20-pc census enables us to identify the nearest star or brown dwarf in each constellation (Table B6). Interestingly, six of the eighty-eight constellations have a Y dwarf as their nearest member despite the fact that Y dwarfs have not yet been fully mapped within this 20-pc volume.  \nNOTE ADDED IN PROOF: The nearby brown dwarf candidate CWISE J165909.91 -351108.5 from Table 2 has been confirmed as a late-L dwarf by Robbins et al. (2023a), but it likely falls outside the 20-pc census. Additionally, Robbins et al. (2023b) find that CWISE J105512.11+544328.3, from Table 14, is not a T subdwarf but rather an anomalously blue Y dwarf.  \nACKNOWLEDGMENTS: Davy Kirkpatrick, Federico Marocco, and Chris Gelino acknowledge support from grant #80NSSC20K0452 awarded for proposal 18-2ADAP180175 under the NASA Astrophysics Data Analysis Program. Data presented here are based on observations obtained at the Hale Telescope, Palomar Observatory as part of a continuing collaboration between the California Institute of Technology, NASA/JPL, Yale University, and the National Astronomical Observatories of China. Some of the data presented herein were obtained at the W. M. Keck Observatory, which is 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 indigenous Hawaiian community. We are most fortunate to have the opportunity to conduct observations from this mountain. This research has made use of the Keck Observatory Archive (KOA), which is operated by the W. M. Keck Observatory and the NASA Exoplanet Science Institute (NExScI), under contract with the National Aeronautics and Space Administration. The first author would like to thank Patrick Shopbell at Caltech for resurrecting an Exabyte drive that successfully read raw CTIO data from 1997. Part of this research was carried out at the Jet Propulsion Laboratory, California Institute of Technology, under a contract with the National Aeronautics and Space Administration (80NM0018D0004). He would also like to thank Mike Cushing for discussion of evolved star loci in the 20-pc colormagnitude diagrams.  \nOne observation reported in this paper was obtained with the Southern African Large Telescope (SALT) under program 2021-2-SCI-027 (PI: Faherty). Roberto Raddi acknowledges support from Grant RYC2021-030837-I funded by MCIN/AEI/ 10.13039/501100011033 and by \'European Union NextGenerationEU/PRTR\', and partial support by the AGAUR/Generalitat de Catalunya grant SGR-386/2021 and by the Spanish MINECO grant PID2020-117252GBI00. Eileen Gonzales acknowledges support from the Heising-Simons Foundation through a 51 Pegasi b Fellowship. This publication makes use of VOSA, developed under the Spanish Virtual Observatory (https://svo.cab.inta-csic. es) project funded by MCIN/AEI/10.13039/501100011033/ through grant PID2020-112949GB-I00. This research made use of the Montreal Open Clusters and Associations (MOCA) database, operated at the Montréal Planétarium (J. Gagné et al., in preparation). Backyard Worlds research was supported by NASA grant 2017-ADAP17-0067 and by the NSF under grants AST- 2007068, AST-2009177, and AST2009136. Johanna Vos acknowledges support from a Royal Society - Science Foundation Ireland University Research Fellowship (URF \\ 1 \\ 221932).  \nFacilities: WISE, Gaia, Spitzer (IRAC), Hale (WIRC, DBSP, TSpec), Keck:I (MOSFIRE, NIRES), IRTF (SpeX), Gemini:South (FLAMINGOS-2), Blanco (RCSPec, ARCoIRIS), NTT (SOFI), Magellan:Baade (FIRE), Shane (Kast), SALT (RSS), SOAR (OSIRIS), ARC (TSpec).  \nSoftware: WiseView (Caselden et al. 2018), Spextool (Cushing et al. 2004; Vacca et al. 2003), MOPEX/APEX (Makovoz & Khan 2005; Makovoz & Marleau 2005), IRAF (Tody 1986, 1993), kastredux (Burgasser et al., in prep.), FIREHOSE/MASE (Bochanski et al. 2009), crowdsource (Schlafly et al. 2018).', 'REFERENCES': "Abazajian, K. N., Adelman-McCarthy, J. K., Agüeros, M. A., et al. 2009, ApJS, 182, 543. doi:10.1088/0067-0049/182/2/543  \nAberasturi, M., Burgasser, A. J., Mora, A., et al. 2014, AJ, 148, 129. doi:10.1088/0004-6256/148/6/129 Aberasturi, M., Caballero, J. A., Montesinos, B., et al. 2014, AJ, 148, 36. doi:10.1088/0004-6256/148/2/36 Abt, H. A. 2017, PASP, 129, 044201. doi:10.1088/1538-3873/aa5b18 Abt, H. A. & Willmarth, D. 2006, ApJS, 162, 207. doi:10.1086/498095 Abt, H. A. 1981, ApJS, 45, 437. doi:10.1086/190719 Abt, H. A. & Levy, S. G. 1976, ApJS, 30, 273. doi:10.1086/190363 Abt, H. A. 1970, ApJS, 19, 387. doi:10.1086/190215 Abt, H. A. 1965, ApJS, 11, 429. doi:10.1086/190120 Adams, W. S., Joy, A. H., Humason, M. L., et al. 1935, ApJ, 81, 187. doi:10.1086/143628 Aerts, C. 2021, Reviews of Modern Physics, 93, 015001.  \ndoi:10.1103/RevModPhys.93.015001  \nAffer, L., Micela, G., Damasso, M., et al. 2016, A&A, 593, A117. doi:10.1051/0004-6361/201628690 Affer, L., Micela, G., Morel, T., et al. 2005, A&A, 433, 647. doi:10.1051/0004-6361:20041308 Agati, J.-L., Bonneau, D., Jorissen, A., et al. 2015, A&A, 574, A6. doi:10.1051/0004-6361/201323056 Akeson, R., Beichman, C., Kervella, P., et al. 2021, AJ, 162, 14. doi:10.3847/1538-3881/abfaff Albert, L., Artigau, É., Delorme, P., et al. 2011, AJ, 141, 203. doi:10.1088/0004-6256/141/6/203 Alberts, F., Savonije, G. J., van den Heuvel, E. P. J., et al. 1996, Nature, 380, 676. doi:10.1038/380676a0 Allen, C., Poveda, A., & Herrera, M. A. 2000, A&A, 356, 529 Allen, P. R., Burgasser, A. J., Faherty, J. K., et al. 2012, AJ, 144, 62. doi:10.1088/0004-6256/144/2/62 Allen, R. H. 1899, New York, Leipzig [etc.] G.E. Stechert, 1899. Allende Prieto, C. & Lambert, D. L. 1999, A&A, 352, 555  \nBenedict, G. F., McArthur, B. E., Franz, O. G., et al. 2001, AJ, 121, 1607. doi:10.1086/319399  \nBenjamin, R. A., Churchwell, E., Babler, B. L., et al. 2003, PASP, 115, 953. doi:10.1086/376696  \nBensby, T., Feltzing, S., & Oey, M. S. 2014, A&A, 562, A71. doi:10.1051/0004-6361/201322631  \nBerdnikov, L. N. & Pastukhova, E. N. 2008, Peremennye Zvezdy, 28, 9  \nBergeron, P., Kilic, M., Blouin, S., et al. 2022, ApJ, 934, 36. doi:10.3847/1538-4357/ac76c7  \nBergeron, P., Dufour, P., Fontaine, G., et al. 2019, ApJ, 876, 67. doi:10.3847/1538-4357/ab153a  \nBergeron, P., Leggett, S. K., & Ruiz, M. T. 2001, ApJS, 133, 413. doi:10.1086/320356  \nBergeron, P., Ruiz, M. T., & Leggett, S. K. 1997, ApJS, 108, 339. doi:10.1086/312955  \nBergeron, P., Ruiz, M.-T., Leggett, S. K., et al. 1994, ApJ, 423, 456. doi:10.1086/173823  \nBergeron, P., Saffer, R. A., & Liebert, J. 1992, ApJ, 394, 228. doi:10.1086/171575  \nBergfors, C., Brandner, W., Bonnefoy, M., et al. 2016, MNRAS, 456, 2576. doi:10.1093/mnras/stv2768  \n- Bergfors, C., Brandner, W., Janson, M., et al. 2010, A&A, 520, A54. doi:10.1051/0004-6361/201014114  \nBernat, D., Bouchez, A. H., Ireland, M., et al. 2010, ApJ, 715, 724. doi:10.1088/0004-637X/715/2/724  \nBernkopf, J., Chini, R., Buda, L.-S., et al. 2012, MNRAS, 425, 1308. doi:10.1111/j.1365-2966.2012.21534.x  \nBerski, F. & Dybczy'nski, P. A. 2016, A&A, 595, L10. doi:10.1051/0004-6361/201629835  \nBessel, F. W. 1844, MNRAS, 6, 136. doi:10.1093/mnras/6.11.136 Best, W. M. J., Liu, M. C., Magnier, E. A., et al. 2021, AJ, 161, 42. doi:10.3847/1538-3881/abc893  \nBest, W. M. J., Liu, M. C., Magnier, E. A., et al. 2020, AJ, 159, 257. doi:10.3847/1538-3881/ab84f4  \nBest, W. M. J., Liu, M. C., Magnier, E. A., et al. 2015, ApJ, 814, 118. doi:10.1088/0004-637X/814/2/118  \nBest, W. M. J., Liu, M. C., Magnier, E. A., et al. 2013, ApJ, 777, 84. doi:10.1088/0004-637X/777/2/84  \nBeuzit, J.-L., Ségransan, D., Forveille, T., et al. 2004, A&A, 425, 997. doi:10.1051/0004-6361:20048006  \nBidelman, W. P. 1980, PASP, 92, 345. doi:10.1086/130674  \nBidelman, W. P. 1985, ApJS, 59, 197. doi:10.1086/191069  \nBihain, G., Scholz, R.-D., Storm, J., et al. 2013, A&A, 557, A43. doi:10.1051/0004-6361/201322141  \nBiller, B. A., Grandjean, A., Messina, S., et al. 2022, A&A, 658, A145. doi:10.1051/0004-6361/202142438  \nBlomme, R., Fremat, Y., Sartoretti, P., et al. 2022, arXiv:2206.05486  \nHoffleit, D., Saladyga, M., & Wlasuk, P. 1983, A Supplement to the Bright Star Catalogue. Containing data compiled through 1981 for stars 7.10 V and brighter that are not in the Bright Star Catalogue.. D. Hoffleit, M. Saladyga, P. Wlasuk.Yale University Observatory,PO Box 6666, New Haven, Conn. 06511, USA. 5+135 pp. Price ∂ 12.00USA, ∂ 13.00 (1983). ISBN 0-914753-01-0.  \nHoffleit, D. & Jaschek, C. 1982, The Bright Star Catalogue. Fourth revised edition. (Containing data compiled through 1979)., by Hoffleit, D.; Jaschek, C.. New Haven, CT (USA): Yale University Observatory, 15 + 472 p.  \nHoffleit, D. 1964, New Haven, Conn.: Yale University  \nObservatory, 1964, 3rd rev.ed., edited by Hoffleit, Dorrit  \nHøg, E., Fabricius, C., Makarov, V. V., et al. 2000, A&A, 355, L27  \nHollands, M. A., Tremblay, P.-E., Gänsicke, B. T., et al. 2018,  \nMNRAS, 480, 3942. doi:10.1093/mnras/sty2057  \nHolmberg, J., Nordström, B., & Andersen, J. 2007, A&A, 475, 519. doi:10.1051/0004-6361:20077221  \nHorch, E. P., Casetti-Dinescu, D. I., Camarata, M. A., et al. 2017, AJ, 153, 212. doi:10.3847/1538-3881/aa6749  \n- Horch, E. P., Bahi, L. A. P., Gaulin, J. R., et al. 2012, AJ, 143, 10. doi:10.1088/0004-6256/143/1/10\n- Horch, E. P., Gomez, S. C., Sherry, W. H., et al. 2011, AJ, 141, 45. doi:10.1088/0004-6256/141/2/45  \nHolberg, J. B., Oswalt, T. D., & Sion, E. M. 2002, ApJ, 571, 512. doi:10.1086/339842  \nHouk, N. 1982, Michigan Catalogue of Two-dimensional Spectral Types for the HD stars. Volume\\_3. Declinations -40\\_ƒ0 to -26\\_ƒ0., by Houk, N.. Ann Arbor, MI(USA): Department of Astronomy, University of Michigan, 12 + 390 p.  \nHouk, N. & Smith-Moore, M. 1988, Michigan Catalogue of  \nTwo-dimensional Spectral Types for the HD Stars. Volume 4,  \nDeclinations -26°.0 to -12°.0.. N. Houk, M.  \nSmith-Moore.Department of Astronomy, University of Michigan, Ann Arbor, MI 48109-1090, USA. 14+505 pp. Price US 25.00 (USA, Canada), US 28.00 (Foreign) (1988).  \nHoward, A. W., Johnson, J. A., Marcy, G. W., et al. 2010, ApJ, 721, 1467. doi:10.1088/0004-637X/721/2/1467  \nHowarth, I. D., Goss, K. J. F., Stevens, I. R., et al. 2014, MNRAS, 440, 1674. doi:10.1093/mnras/stu398  \n- Hsu, C.-C., Burgasser, A. J., Theissen, C. A., et al. 2021, ApJS, 257, 45. doi:10.3847/1538-4365/ac1c7d\n- Huber, D., Bryson, S. T., Haas, M. R., et al. 2016, ApJS, 224, 2. doi:10.3847/0067-0049/224/1/2\n- Huber, D., Matthews, J. M., Croll, B., et al. 2009, A&A, 505, 715. doi:10.1051/0004-6361/200912139  \nHumphreys, A., Meisner, A. M., Burgasser, A. J., et al. 2023, Research Notes of the American Astronomical Society, 7, 184. doi:10.3847/2515-5172/acf4a0", 'APPENDIX': '(FLAMINGOS-2; Eikenberry et al. 2006, Elston et al. 2003, Jannuzi & Bechtold 2004) at the 8.1-m Gemini-South Observatory on Cerro Pachón, Chile, on the nights of 2014 December 01, 2015 June 30, 2019 April 14, 2019 April 29, 2021 February 22, 2021 July 17, 2021 July 21, and 2021 October 18 UT. Photometric acquisition and reductions followed the procedures described in section 9.1 of Meisner et al. (2020a).  \nA.1.4. SOAR/OSIRIS  \nOne object was observed with the Ohio State Infra-Red Imager/Spectrometer (OSIRIS) at the 4.1-m SOuthern Astrophysics Research (SOAR) Telescope on Cerro Pachón, Chile, on the night of 2012 March 12 UT. Photometric acquisition and reduction of this lone H -band data point followed the methodology outlined in section 2.2.5 of Kirkpatrick et al. (2012).', 'A.1.5. Spitzer/IRAC': 'The Spitzer Heritage Archive was queried for directed or serendipitous observations of objects in our 20-pc candidate list (Table 2). The locations of thirty-one of these candidates were found to have Spitzer/IRAC observations in ch1 and/or ch2. The full list is shown in Table A1. Although some of the Spitzer data were too shallow to detect our objects or were obtained at an epoch when our candidate was blended with a background source, most objects had measurable photometry. We used the MOPEX/APEX software (Makovoz & Khan 2005, Makovoz & Marleau 2005) on each Astronomical Observation Request (AOR) to create mosaics, perform source detection, and then measure the photometry using the stack of individual frames. The output of the APEX code is the flux, in µ Jy, for each detection using both aperture and PRFfit techniques. The Spitzer photometry reported in Table 2 is the PRF-fit photometry after converting to magnitudes using the correction factors listed in Table C.1 of the IRAC Handbook 46 along with the flux zeropoints in each band, as given in Table 4.1 of the Handbook. For objects having multiple AORs for a band, the reported photometry in that band is a weighted mean of the individual measurements in each AOR.', 'A.1. Photometry': 'For possible M, L, T, and Y dwarf additions to the 20pc census, we have searched for published near- and midinfrared photometry using online surveys such as the Two Micron All Sky Survey (2MASS; Skrutskie et al. 2006), the VISTA Hemisphere Survey (VHS; McMahon et al. 2013), the UKIRT Hemisphere Survey (UHS; Dye et al. 2018), and the Wide-field Infrared Survey Explorer (WISE; Wright et al. 2010). These photometric measurements are listed in Table 2. In other cases, we have obtained our own ground-based follow-up or have searched the Spitzer Heritage Archive 45 for images with which to measure photometry. These results are also presented in Table 2 but discussed further in the subsections below.', 'A.1.1. Palomar/WIRC': "Eighteen objects were observed with the Wide-field Infrared Camera (WIRC; Wilson et al. 2003b) at the Hale 5-m telescope on Palomar Mountain during the nights 2014 July 03, 2014 September 14, 2016 February 26, 2018 September 01, 2019 July 14, 2020 February 05, 2020 June 03, 2020 July 03, 2020 September 03, 2020 October 09, 2021 July 01, and 2021 August 10 UT. Data were acquired in the Maunakea Observatory filter system's J and H bands. Our standard observing technique, calibration strategy, and reduction methodology have been discussed in section 3.1.5 of Kirkpatrick et al. (2011) with updates as discussed in section 9.2 of Meisner et al. (2020a).", 'A.1.2. Keck/MOSFIRE': "Thirteen objects were observed with the Multi-Object Spectrometer For Infra-Red Exploration (MOSFIRE; McLean et al. 2012) at the 10-m W. M. Keck I telescope on Maunakea, Hawai'i, on the nights of 2021 August 27 and 2022 January 21 UT. Photometric acquisition and reductions followed the procedures described in section 3.1.1 of Schneider et al. (2021).", 'A.1.3. Gemini-South/FLAMINGOS-2': 'Nine objects were observed with the FLoridA Multi-object Imaging Near-infrared Grism Observational Spectrometer 2', 'A.2. Spectroscopy': 'To aid in the characterization of objects, spectroscopy was acquired of 20-pc members discovered by Gaia, 20-pc suspects discovered by the Backyard Worlds: Planet 9 citizen science group, or 20-pc members lacking published spectral types. These reduced spectra are illustrated in Figures A1 through A5. Ten different instruments were used for this follow-up, as detailed below and summarized in Table A2.', 'A.2.1. CTIO/RCSpec': 'Optical follow-up of five objects was obtained on the UT dates of 1995 August 13-14, 1997 July 14, and 1997 July 16 at the Cerro Tololo Interamerican Observatory (CTIO) 4m telescope and on 1996 May 20 at the CTIO 1.5m telescope using the R-C Spectrograph with Folded Schmidt Camera (RCSpec). For the 1995 and 1996 observations, a 300 line mm -1 grating with a GG 495 order-blocking filter was used with the 1024 × 1024 CCD to cover a wavelength range from 6050 to 9550 Å. For the 1997 observations, a 316 line mm -1 grating with an OG 515 order-blocking filter was used with the Loral 3K CCD to cover a useable wavelength range from 5200 to 10000 Å. In addition to the targets, standard calibrations - biases, dome flats, arcs, and flux calibration standards - were also obtained. Reductions were accomplished using the Image Reduction and Analysis Facility ( IRAF ; Tody 1986, 1993), as described in Kirkpatrick et al. (1997).', 'A.2.2. Lick/Kast': 'The Kast Double Spectrograph at the Lick 3m Shane Telescope was used for optical follow-up of six objects on UT dates 2019 September 20, 2020 March 06, 2020 August 1516, 2020 December 14, and 2022 July 02. The only data used were from the red arm, which employed a 600 line mm -1 grating blazed at 7500 Å to cover the wavelength range from ∼ 6000 to ∼ 9000 Å. In addition to standard wavelength and flux calibrations, G2V and A0V stars were obtained near in time and on sky to the targets to correct for telluric absorption. The data were reduced using the kastredux 47 package (Burgasser et al., in prep.), as further described in Schneider et al. (2021).', 'A.2.3. Palomar/DBSP': 'The Double Spectrograph (DBSP; Oke & Gunn 1982) at the Hale 5m telescope on Palomar Mountain was used for sixty additional follow-up spectra. The UT nights of observation were 1995 December 02, 2021 December 06, 2022 January 02, 2022 February 03, 2022 February 07, 2022 May 24, 2022 May 30, 2022 June 07, and 2022 August 27. For the 1995 run, the D68 dichroic was used to split the light at ∼ 6800 Å between the two arms. Gratings with 300 line mm -1 blazed at 3990 Å and with 316 line mm -1 blazed at 7150 Å were used in the blue and red arms, respectively, producing continuous wavelength coverage from ∼ 5100 to  \n9200 Å. For the 2021 and 2022 runs, the D55 dichroic was used instead to split the light near 5500 Å. A 600 line mm -1 grating blazed at 3780 Å was used in the blue arm and a 316 line mm -1 blazed at 7150 Å was used in the red arm, producing coverage from ∼ 3300 to 5500 Å on the blue side and from ∼ 5700 to 10000 Å on red side. (For the 2022 June run, fringing in the blue arm caused data shortward of 4500 Å to be unusable.) Standard calibrations and IRAF data reductions were employed, as further described in Kirkpatrick et al. (1991).', 'A.2.4. SALT/RSS': "An additional optical spectrum was acquired with the Robert Stobie Spectrograph (RSS; Burgh et al. 2003; Kobulnicky et al. 2003) on the 11.1 × 9.8-m Southern African Large Telescope (SALT; Buckley et al. 2006) on 2021 December 26 UT. The spectrograph was used in long slit mode using the PG0900 grating at an angle of 20 · , which produces coverage over the ranges 6033-7028, 7079-8045, and 8091-9023 Å across the 3 × 1 mini-mosaic. Our reductions began with the observatory-provided pre-processed data, for which gain correction, correction for cross-talk, and overscan subtraction had been applied. We then wavelength calibrated using neon arc lines obtained immediately after the target's spectroscopic data and flux calibrated using the Hamuy et al. (1994) standard EG21 acquired with the same spectroscopic setup on 2023 January 24 UT.", 'A.2.5. APO/TSpec': 'TripleSpec (Wilson et al. 2004) on the ARC 3.5m telescope at Apache Point Observatory was used for near-infrared follow-up of two objects on the nights of 2018 September 23 and 2019 October 08 UT. The spectrograph provides spectral coverage from 0.95 to 2.46 µ m across five spectral orders. Data were taken with the conventional near-infrared technique of nodded pairs to perform background/bias subtraction, and standard calibrations were also acquired, including quartz lamps for flat fielding and A0 stars for telluric correction and flux calibration. Wavelength calibration was accomplished using night sky lines. Data reduction used Tspectool 48 , a modified version of Spextool (Cushing et al. 2004) rewritten specifically for APO/TripleSpec.', 'A.2.6. CTIO/ARCoIRIS': "Three objects were observed on the nights of 2018 April 02-03 UT using the Astronomy Research with the Cornell Infra Red Imaging Spectrograph (ARCoIRIS) at the 4m Victor Blanco telescope at CTIO. Spectra are acquired across six cross-dispersed orders covering 0.8 to 2.4 µ mat a resolving power of ∼ 3500. Science exposures were taken with AB nod positions along the slit, which has a fixed width of 1 . '' 1. Standard calibrations were acquired as discussed in Greco et al. (2019), and data were reduced using a modified version of  \nFigure A1. Optical spectroscopic follow-up for objects classified as early-M through mid-M. Each target object (black) is normalized to one at 7500 Å and overplotted (in other colors) with the spectral standard nearest the same spectral type. Integral offsets have been added to separate the spectra vertically. Target objects are labeled with brief RA/Dec (hhmm ± ddmm) identifiers. The two target objects at upper left 1839+0901 and 0623+1018 - have been smoothed to improve the signal-to-noise in each wavelength bin. Most spectra have not been corrected for earth's atmospheric absorption, so the contaminating B- and A-bands of O2 at ∼ 6850-6900 and ∼ 7600-7700 Å and telluric bands of H2O at ∼ 7150-7300, 8150-8350, and 9000-9600 Å remain.  \n<!-- image -->  \nthe Spextool package (Cushing et al. 2004), which utilizes A0 stars for telluric correction and flux calibration following the methodology of Vacca et al. (2003).", 'A.2.7. IRTF/SpeX': 'The SpeX instrument on the NASA Infrared Telescope Facility (IRTF) was used for near-infrared spectroscopy of twenty-seven objects over the nights of 2018 June 16, 2018 November 25, 2019 January 23, 2019 March 16-17, 2020 October 30, 2020 November 25, 2021 May 31, 2021 June 30, 2021 September 11, 2021 October 23, 2022 January 09, 2022 January 19, 2022 February 12, 2022 February 21, 2022 March 07, and 2022 March 11 UT. Two different setups were employed. The first, used mainly for brighter targets, was a cross-dispersed mode that provides spectra over the range 0.9-2.4 µ mat a resolving power of R ≡ λ/ ∆ λ ≈ 1200. The second, used primarily for fainter targets, was the prism mode that provides spectra over the range 0.8-2.5 µ mat a resolving power of R ≡ λ/ ∆ λ ≈ 100 -150. As discussed in the subsections above, standard near-infrared calibrations were  \nobtained, and data were reduced using Spextool (Cushing et al. 2004, Vacca et al. 2003).', 'A.2.8. Keck/NIRES': 'Three objects were observed over the nights of 2019 February 14, 2020 July 07, and 2021 February 24 UT using the Near-Infrared Echellette Spectrometer (NIRES; e.g., Wilson et al. 2004) at the 10m W. M. Keck II telescope. These data provided spectral coverage from 0.94 to 2.45 µ m. Setup and calibrations were identical to those described in Meisner et al. (2020b), and reductions used a modified version of Spextool (Cushing et al. 2004, Vacca et al. 2003).', 'A.2.9. Magellan/FIRE': 'The Folded-port Infrared Echellette spectrograph (FIRE; Simcoe et al. 2013) at the 6.5 m Walter Baade (Magellan I) telescope at Las Campanas Observatory was used to observe three objects over the nights of 2016 January 23, 2019 December 11, and 2020 February 18 UT. Observations were done in prism mode, which covers the range from 0.80 to  \nFigure A2. Optical spectroscopic follow-up for objects classified as mid-M through early-L. Each target object (black) in the two left panels is normalized to one at 7500 Å and overplotted (in other colors) with the spectral standard nearest the same spectral type. In the far right panel, this normalization is done instead at 8250 Å. Offsets in steps of 1.5 have been added to separate the spectra vertically. A few target objects - 1921 -2915 (M7), 1906+4011 (L1), 0033+4340 (L2), and 0617+1945A (L2) - have been smoothed to improve the signal-to-noise in each wavelength bin. The spectrum of 1921 -2915 (M7) also suffers from residual cosmic ray hits. See the caption to Figure A1 for more details.  \n<!-- image -->  \n2.45 µ m. Standard calibrations were acquired, and data were reduced using the FIREHOSE pipeline, which is based on the MASE (Bochanski et al. 2009) and Spextool (Cushing et al. 2004, Vacca et al. 2003) packages.', 'A.2.10. Palomar/TSpec': "Finally, seven near-infrared spectra were acquired with the Triple Spectrograph (TSpec; Herter et al. 2008) at Palomar Mountain's 5m Hale telescope on the nights of 2018 April 28-29, 2018 October 17, and 2019 September 18 UT. Setup and calibrations were identical to those described in Kirkpatrick et al. 2011. As with many of the other near-infrared spectroscopic data sets discussed above, TSpec data were also reduced with a modified version of Spextool (Cushing et al. 2004, Vacca et al. 2003).", 'A.2.11. Analysis': 'Spectral classification was accomplished by comparing spectra of the target objects to established on-sky anchors for each integral spectral type. For the optical spectra, these anchor points were taken from Kirkpatrick et al. (1991) for objects of type mid-K through late-M and from Kirkpatrick et  \nal. (1999) for L dwarfs (Figures A1-A2). Optical classifications of objects earlier than type K (Figure A3) used spectral anchors taken from Gray & Corbally (2009). Near-infrared classification (Figures A4-A5) for M dwarfs, L dwarfs, and early-T dwarfs used the anchors described in Kirkpatrick et al. (2010), with the rest of the T dwarf anchors coming from Burgasser et al. (2006). For more on the methodology employed for both optical and near-infrared classifications, see Kirkpatrick et al. (2010).  \nThese classification anchors are generally old disk objects with metallicities similar to the Sun. In a few cases, described below, the target object failed to match an anchor spectrum because of anomalous features attributable to extreme youth, lower metallicity, or other reasons including unresolved binarity. These special cases are addressed further below:  \nCWISE J045334.34+203350.2: At J -band, this object best matches the L5 standard, but there are clear discrepancies with the L5 standard across all wavelengths (Figure A4). The continuum of 0453+2033 is much flatter between 1.1 and 1.3 µ m, the FeH band at 0.99 µ m is much stronger, and the H - and K -band portions emit less flux rel-  \nFigure A3. Optical spectroscopic follow-up of objects not classified as M dwarfs or L dwarfs. Each target object is normalized to one at its peak flux. Objects in the far left panel are hot stars, and objects in the two right panels are colder stars or other background objects. Integral offsets have been added to separate the spectra vertically. A few spectra - 0213 -3345 (wd), 1955+2224 (mid-F?), 1318+3810 (early-G), and all of those in the rightmost panel - have been smoothed to improve the signal-to-noise in each wavelength bin. See the caption to Figure A1 for more details.  \n<!-- image -->  \nive to J -band than does the standard. We find that the J -band spectrum of 0453+2033 is a better match to 2MASS J17561080+2815238, which is typed in both the optical and the near-infrared as an sdL1 (Kirkpatrick et al. 2010, Zhang et al. 2018), in both the continuum shape and the strength of the FeH band. However, 0453+2033 has more flux at H and K relative to J than does 2MASS J1756+2815, possibly indicating that the former is a slightly later subdwarf. Given that the set of anchors for the L subdwarf spectral sequence is still incomplete (Zhang et al. 2018), we tentatively classify this object as an early- to mid-sdL.  \nCWISE J055942.94 -012002.4: Of the spectra in Figures A4-A5 that have a "red" or "slightly red" classification, only 0559 -0120 has the triangular-shaped H -band peak indicative of low gravity. Such low-gravity objects are necessarily young, as they have not yet contracted to their final radii. Using just the sky position and proper motion values (Table 2), as its parallax and radial velocity have not yet been measured, BANYAN Σ (Bayesian Analysis for Nearby Young Associations; Gagné et al. 2018) gives the object an  \n80% chance of belonging to a known, young moving group - either the AB Doradus group or, less likely, the β Pictoris group. If an AB Dor member, BANYAN Σ predicts 46 ± 3 pc with a radial velocity of 22 ± 2 km s -1 ; if a β Pic member, the predictions are d = 21 ± 3 pc and RV = 19 ± 2 km s -1 . Using solely an MKs vs. spectral type relation (Dupuy & Liu 2012), as advocated for young L dwarfs in Schneider et al. (2023), we estimate a distance of ∼ 28.8 pc for this L5.5 dwarf, based on a value of Ks =14 . 34 ± 0 . 09 mag from the 2MASS All-Sky Point Source Catalog.  \nCWISE J075227.38+053802.6: We classify this object as L9 pec (Figure A5). The peculiarities stem from the two unusual absorption troughs at 1.63 and 1.67 µ m within the H -band plateau and the unusual inflection near 2.21 µ m at K -band. Such features are indicative of methane absorption, which should not be present shortward of 2.5 µ m in an L9 dwarf. As previous papers such as Burgasser (2007b), Burgasser et al. (2010b), and Bardalez Gagliuffi et al. (2014) have noted, such spectra may indicate the presence of an unresolved binary comprised of two morphologically distinct  \nFigure A4. Near-infrared spectroscopic follow-up of objects classified as mid-M through mid-L. Each target object (black) is normalized to one at 1.28 µ m and overplotted (in other colors) with the spectral standard nearest the same spectral type. Integral offsets have been added to separate the spectra vertically. Target objects are labeled with brief RA/Dec (hhmm ± ddmm) identifiers. One spectrum - 0907 -4308 (L5) - has been smoothed to improve the signal-to-noise in each wavelength bin. Data deep within the telluric water bands near ∼ 1.4 and ∼ 1.75 µ m are not displayed for some targets because of the poor signal-to-noise in those regions.  \n<!-- image -->  \nspectra - a non-methane M or L dwarf and a methane-rich T dwarf. If 0752+0538 represents such an unresolved binary, modeling (see section 4.5 of Kirkpatrick et al. 2016) suggests it is likely a late-L plus early-T composite system.  \nCWISE J075853.12 -232645.8: We classify this object as T2.5 pec (Figure A5). Although its J -band peak matches both the T2 and T3 standards equally well, the H -band flux is suppressed and the K -band flux elevated relative to the standards. We find that a synthetic spectrum made up of components of types L8-L9 and T5-T6 fits the overall spectral shape slightly better than the single standards, suggesting perhaps that this object is an unresolved binary.  \nCWISE J132403.81 -052631.4: The width of the J -band peak in this object best matches that of the T4 standard (Figure A5), but the fits at both shorter and longer wavelengths are much poorer. Specifically, the H - and K -band portions of 1324 -0526 emit less flux relative to J -band than does the standard, and the K -band portion is notably flattened, an effect often ascribed to increased collision-induced absorption by H2. Moreover, the Y -band portion emits more flux relative to J -band than does the standard. Elevated Y -band flux  \nand suppressed K -band flux are seen in a comparison of the sdT5.5 dwarf HIP 73786B (Figure 1 of Zhang et al. 2019) to standards of type T5 and T6, although the discrepancies are stronger in 1324 -0526, and the latter also shows suppressed H -band flux. In the case of HIP 73786, the system has a measured subsolar metallicity of [Fe/H] = -0 . 3 ± 0 . 1 (Murray et al. 2011) from the K5 V primary, suggesting that the metallicity of 1324 -0526 is somewhat lower still. We classify 1324 -0526 as sdT4.  \nCWISE J221859.41+114642.7: The width of this object\'s J -band peak is most similar to the T7 standard, but its H - and K -band flux peaks are suppressed, with the latter being noticeably flattened. As with 1324 -0526 above, such features are typical of subdwarfs, although the suppression of the Y -band peak in 2218+1146 runs contrary to the trend seen in T subdwarfs of slightly earlier type. In the sdT8 WISE J200520.38+542433.9, a wide companion in the low-metallicity ([Fe/H] = -0 . 64 ± 0 . 17) Wolf 1130 system (Mace et al. 2013b), the Y -band peak is shifted notably to the blue - from 1.09 to 1.03 µ m - relative to the standards, an effect also seen in the isolated sdT6.5 dwarf ULAS  \nFigure A5. Near-infrared spectroscopic follow-up of objects classified as mid-L through late-T. One spectrum - 0133+8031 (T4) - has been smoothed to improve the signal-to-noise in each wavelength bin. Data deep within the telluric water bands near ∼ 1.15, ∼ 1.4, and ∼ 1.75 µ m are not displayed for some targets because of the poor signal-to-noise in those regions. See the caption to Figure A4 for other details.  \n<!-- image -->  \nJ131610.28+075553.0 (Burningham et al. 2014). Our spectrum is too noisy in this region to determine whether the same  \neffect is present in 2218+1146, so we classify this object as T7 pec pending further confirmation of its subdwarf status.  \nTable A2 . Spectroscopic Follow-upTable A2 continued  \nTable A2 (continued)  \nTable A2 continued  \nTable A2 (continued)  \nTable A2 (continued)', 'A.3. Astrometry': 'Additional parallaxes have been measured as part of an ongoing ground-based program and through serendipitous imaging data found in the Spitzer Heritage Archive. These results are discussed further below.', 'A.3.1. NPARSEC results': 'Nearby objects continue to be targeted as part of the NTT PARallaxes of Southern Extremely Cool objects (NPARSEC) project, a long-term program (186.C-0756 with R. Smart, PI; 105.C-0781, 108.21XQ.0001, and 108.21XQ.002 with E. Costa, PI) using the infrared spectrograph and imaging camera Son OF ISAAC (SOFI; Moorwood et al. 1998) on the New Technology Telescope (NTT). The observational methodology and reduction procedures are identical to those  \ndiscussed in Smart et al. (2013). For the eleven objects listed in Table A3, the new NPARSEC preliminary values have smaller uncertainties than previously published parallaxes. The table gives the object names, J2000 coordinates, mean epoch of observation, the absolute parallax, the correction applied to the relative parallax to convert to absolute, the proper motion values in Right Ascension and Declination, the total time baseline of the NTT observations, the number of reference stars used, and the total number of separate observational epochs.  \nNine of the targets have absolute parallaxes greater than 50 mas, but for SDSS J163022.92+081822.0 and 2MASS J23312378 -4718274, these better determined parallaxes have values below 50 mas, so we now exclude these three from the 20-pc census. We note, however, that these results are still considered preliminary and will be finalized once the NPARSEC program draws to a close.  \nTable A3 . Preliminary Parallax and Motion Fits for Objects on the NPARSEC Parallax Programs  \nTable A3 (continued)', 'A.3.2. Spitzer results': "CWISE J181125.34+665806.4 (hereafter 1811+6658; see Figure A6) is located only 1.2 degrees from the north ecliptic pole (NEP), and the area around the NEP was routinely observed by the Spitzer Space Telescope. As shown in Table A1, the location of 1811+6658 was observed repeatedly in post-cryogenic programs 10147 (PI: Bock) and 13153 (PI: Capak) in an attempt to explore the genesis of fluctuations in the extragalactic background light and to provide IRAC/ch1 and IRAC/ch2 data on touchstone fields that will be used by Euclid, Roman, and JWST to study galaxy growth during the epoch of reionization. The data in program 10147 cover the timeframe from May 2014 to Sep 2014, and those in program 13153 cover Feb 2017 to Feb 2019.  \nTo extract astrometry from these data sets, we searched for blocks of ch2 coverage that had sufficient depth and redundancy to provide a similar per-epoch astrometric accuracy to that obtained in our own Spitzer parallax programs (Kirkpatrick et al. 2019, 2021a). (This cold brown dwarf is much brighter at ch2, 15.95 mag, than at ch1, 18.23 mag, so only the longer-wavelength band would provide sufficient signal-to-noise for our astrometric needs.) Program 10147 used 30s exposures per frame, and the position of 1811+6658 was observed at four or fewer epochs. The data from program 13153, on the other hand, used 100s exposures per frame and had more coverage at each sky position.  \nWe pared this data set down to include only those frames for which the location of 1811+6658 was far enough from the frame edge to provide a reasonable number of Gaia DR3 reference stars surrounding the target's location. Specifically, we retained only those ch2 frames that imaged all six of our pre-selected Gaia DR3 astrometric reference objects encircling a 60 '' zone centered on the location of 1811+6658.  \nFor program 10147, this left only two or three frames per epoch; this lack of redundancy combined with the short exposure time means that these data are unsuitable for astrometric analysis. For program 13153, however, we are left with four to six redundant, longer exposure frames per epoch (defined here to be per AOR), which is suitable for our reduction techniques. Of those program 13153 AORs listed in Table A1, only the ch2 data in 62377728, 65133312, 68615680, and 68631296 lacked sufficient redundancy. The time span covered by the remaining data sets is Jul 2017 to Jan 2019. These data were extracted and astrometrically calibrated to the Gaia DR3 reference frame as described in (Kirkpatrick et al. 2021a).  \nGiven that the usable Spitzer data only cover a year and a half, we turned to WISE astrometry to provide the additional baseline needed to disentangle parallax from proper motion. The NEP is within the boresight of the WISE spacecraft on every orbit, but given the 47 ' -wide field of view, 1811+6658 is not within the continuous viewing zone. However, that location is viewed by WISE during a span of 50+ days every six months as the scan pattern rotates around the ecliptic pole. As such, there are several weeks of coverage twice per year covering its location.  \nWe ran the crowdsource detection software (Schlafly et al. 2018) on time-resolved unWISE coadds (Meisner et al. 2018) for all ten-day epochal mosaics covering the position of 1811+6658. We retained those source lists for which the frame coverage depth at the location of 1811+6658 was 40 or greater. This was done in an effort to assure that the area surrounding the target's location also had sufficient coverage, as this area is needed for the astrometric calibrators. The measured positions of these surrounding astrometric standards were used to place the measured position of 1811+6658 onto the same Gaia DR3 reference frame used for the Spitzer observations.  \nFigure A6. Cutout images, 120 '' on a side with north up and east to left, of 1811+6658. (Left) WISE data at epoch 2021.6. The separate W1 and W2 bands have been mapped into a color scheme in which objects appearing at roughly equal brightness in each will appear black, and those appearing primarily in W2 will appear orange (Caselden et al. 2018). The brown dwarf 1811+6658 is the orange object at the center of the field. (Center) PanSTARRS data. Bands y / i / g have been mapped into red/green/blue. Note the two blue background sources lying near the center of the field, which is shown in the zoomed inset. (Right) Keck/MOSFIRE data at epoch 2021.7. The detection of 1811+6658 is marked with a red circle (matched to the size of the aperture used in our photometric reductions) and is sandwiched between the two blue background sources seen in the PanSTARRS view. The inset shows a zoom of the field center.  \n<!-- image -->  \nThis astrometry from Spitzer and WISE is listed in Table A4. A fit to the parallax and proper motion was performed on the combined astrometry using the methodology outlined in Kirkpatrick et al. (2021a), resulting in the values shown in Table A5. The results of this fit are shown graphically in Figure A7. We find that the object has a distance of 14.3 + 1 . 6 -1 . 2 pc and a value of Mch 2 = 15 . 16 ± 0 . 21 mag. A comparison to Figure 16d of Kirkpatrick et al. (2021a) suggests a spectral type of early-Y for this absolute magnitude. We further note that the measured colors JMKO -ch2 = 5.66 ± 0.04 mag, W1 -W2 = 3.04 ± 0.09 mag, and ch1 -ch2 = 2.28 ± 0.02 mag - suggest a slightly earlier spectral type of around T9T9.5 based on Figures 16e, 16g, and 16h of Kirkpatrick et al. (2021a). A comparison of our Keck/MOSFIRE JMKO , and WISE W1+W2 images with data from PanSTARRS (Figure A6) shows that 1811+6658 is passing near two blue PanSTARRS sources. Given the low spatial resolution of the WISE (and Spitzer) data, we believe that our measurements of W1 (and ch1) are contaminated by these background objects. The higher resolution of the Keck/MOSFIRE data allows us to separate all three components, but our aperture photometry at JMKO is likely still compromised given our aperture radius of 6 pixels (1 . '' 1). Therefore, our measured JMKO -ch2, W1 -W2, and ch1 -ch2 color are all likely bluer than their true values, supporting our assertion of a Y dwarf spectral type.", 'B. THE LIST OF PROXIMAL SYSTEMS': 'Despite recent WISE-based discoveries of the L+T dwarf binary system WISE J104915.57 -531906.1 AB (Luhman 2013; 1.99 pc distant) and the Y dwarf WISE J085510.83 -071442.5 (Luhman 2014; 2.28 pc distant) adding to our knowledge of the Sun\'s immediate neighbors, Proxima Centauri (Innes 1915; 1.30 pc distant) remains the closest. It is often just referred to as "Proxima", Latin for "nearest", because it is the nearest to the Sun. Yet, its full name translates to "Nearest of Centaurus". This has led some curious individuals to wonder what are the nearest stars - i.e., the other proximal objects - of each of the other eighty-seven official constellations.  \nTable 4 allows us to answer this question, given our current knowledge of the 20-pc census. Proxima Centauri and its primary, α Centauri AB, represent a rare multi-object system for which the parallaxes of the individual components are so accurate that we can determine the far-flung companion to be closer to us than its host binary. For other multiobject systems, discerning the closest component may be far more difficult. Using a short-period binary as an example, determining the closest object in the system depends upon the orbital period and orientation with respect to the Sun, as one component may be the closer one at some epochs and the more distant one at others. Hence, we will identify only the proximal systems in each constellation when a multi-object system arises.  \nTable B6 lists these proximal systems. As examples, the closest system in Canis Major is the binary Sirius AB, and the closest in Delphinus is the Y dwarf WISEPC J205628.90+145953.3. For several constellations, the proximal system is still ambiguous, given current uncertainties in the measured trigonometric parallaxes of the closest can-', 'CWISE J181125.34+665806.4 (B = 88.809)': '<!-- image -->  \n<!-- image -->  \nFigure A7. Best fit to the parallax and proper motion of 1811+6658. (Upper left) Sky plot showing the track of the object along the sky. Black points with large uncertainties are the 57 individual unWISE time-resolved measurements. The orange curve shows the best fit as seen from the vantage point of WISE, and the blue curve shows this same fit from the vantage point of Spitzer. Red lines connect each observation to its predicted point along the best-fit curve. (Upper right) The parallax solution (green) with the proper motion component subtracted out. For clarity, only the 10 Spitzer data points are shown. Red lines connect the times of the Spitzer observations to their predicted points on the curve. (Lower left) The parallactic fit (green) as a function of time in RA and Dec, along with the measured Spitzer astrometry. (Lower right) Residuals around the parallactic fit as function of time in both RA and Dec. Blue lines mark residuals of zero. For additional info on this plot, see Figure 2 of Kirkpatrick et al. (2021a).  \n<!-- image -->  \n<!-- image -->  \nTable A4. Astrometry on the Gaia DR3 Reference Frame for 1811+6658  \nNOTE-(This table is available in its entirety in a machine-readable form in the online journal. A portion is shown here for guidance regarding its form and content.)  \nTable B6 (continued)  \nTable B6 continued  \nTable A5. Parallax and Motion Fit for 1811+6658  \nNOTE-The RA and Dec values are listed on the ICRS coordinate system. The last three rows represent the number of Spitzer ch2 epochs (# Spitzer ) and the number of unWISE W2 epochs (# WISE ) used in the fits, along with the number of five-parameter Gaia DR3 stars used for the astrometric reregistration (# Gaia ).  \ndidates. Constellations having two objects within 1 σ of the same closest value are indicated by footnotes. Note that all constellations have a proximal system within the 20-pc limit of Table 4, the most distant being the K dwarf HD 200779 at 15.05 pc, the closest known object in Equuleus.  \nTable B6 . The Proximal Systems for Each Constellation  \nTable B6 continued  \nTable B6 (continued)  \nIt is worth noting the prevalence of brown dwarfs in Table B6. There are four T dwarf companions residing in proximal systems with K or M dwarf primaries, one L+T binary as its own proximal system, and eleven solivagant L, T, or Y dwarfs that are proximal objects in their own right. In this latter group, over half (six) of the these are Y dwarfs, a spectral type that is not yet fully sampled near the Sun.'}