662B . THE ASTRONOMICAL JOURNAL VOLUME 101, NUMBER 2 FEBRUARY 1991 101. THE LATE-M DWARFS M. S. Bessell Mount Stromlo and Siding Spring Observatories, Institute of Advanced Studies, The Australian National University, Private Bag, 1991AJ Weston Creek Post Office, ACT 2611 Canberra, Australia Received 26 June 1990; revised 30 October 1990 ABSTRACT Far-red spectra and VRIJHK photometry have been obtained for a sample of late-M dwarfs selected on the basis of large reduced red magnitudes from the LHS Catalogue. Half of the stars in the three faintest 1 mag bins are late-M stars, the other red stars are metallic-hydride subdwarfs. Relations between various colors for the late-M dwarfs are investigated. Of all the colors I — K most reliably correlates with spectral type. FeH bands near 9900 A are clearly seen in the spectra of all dwarf stars later than M5. Two stars cooler than VB10, and similar in temperature to LHS2924 have been identified; both have Ha in emission and appear variable in magnitude and R — I color; one is a flare star. The other stars are of earlier spectral type and resemble W359 and VB8. The observed Mj, I — K main sequence is in good agreement with the IG theoretical main sequence of Stringfellow, and the faintest stars could be —0.09 red dwarfs or lower mass brown dwarfs. 1. INTRODUCTION and approximately one half of the observed stars in the top 3 mag bins have been found to be late-M stars with R — / col- Low luminosity M dwarfs are of interest for many rea- sons, both observational and theoretical. The faintest M ors redder than 2.0, the remainder were mainly subdwarf M dwarfs will include stars undergoing gravitational collapse stars with R — / colors between 0.6 and 1.2 and having spec- to the Main Sequence, stars burning hydrogen on the Main tra dominated by bands of CaH (see e.g., Bessell 1982). Sequence, and stars with too low a mass to burn hydrogen, These subdwarfs will be discussed elsewhere (Bessell & Lie- the “brown dwarfs.” It is important to establish whether we ber! 1991). A few stars had continuous spectra and are as- can distinguish between these groups of stars observational- sumed to be white dwarfs. In this paper we will discuss the ly, and whether evolutionary modeling can predict differ- photometric and spectroscopic observations of the late-M ences; however, comparison between theory and observa- dwarfs and review the state of knowledge of the lower main tions is insecure because realistic model atmospheres are not sequence. yet available for M dwarfs. Mould (1976) computed model atmospheres for M dwarfs but these are inadequate at the 2. OBSERVATIONS lowest temperatures because the H2 O opacity was too large. Allard ( 1990) and Kui ( 1990) are computing better atmo- Many people have made observations of the spectra and spheres, but these will necessarily not include the effects of colors of M dwarfs, and we have made use of much of this “star spots” nor non-LTE phenomena such as flares, which data to place the properties of the newly identified M dwarfs can considerably influence the surface layers in the late-M in context. The work of Eggen (e.g., Eggen 1968) and Weis- dwarfs. trop ( 1976), on the R/colors of M stars showed that reliable A preeminent reason for interest in the M dwarfs is the distances were derived from R — I photometric parallaxes, fact that the “missing mass” in the solar neighborhood could and that these were more reliable than the distances derived reside in the faint red and brown dwarfs. Several star-count from the objective prism spectral types, as done by Murray «fe analyses (e.g., Luyten 1968; Gilmore & Reid 1983) have Sanduleak ( 1972). Most recent M dwarf surveys have there- been made which conclude that the observed luminosity fore been undertaken using photometry, either photograph- function turns over at M4 stars, and that the cooler and ic or photoelectric, in the F, R, or / bands. fainter M stars cannot contribute enough mass. Hawkins FR/photometry has been published by Weistrop (1976, (1986) and Hawkins & Bessell (1988), however, by pushing 1979,1980), Eggen (1979,1980), Cousins ( 1980), Théétf a/. hypersensitized IVN and IIIaF plates to their limits, have (1981), Weis (1984, 1986, 1987) and Bessell (1983, 1990). found larger numbers of redder stars, which raises again the Transformations between these related photoelectric sys- possibility that the faintest M dwarfs may contribute signifi- tems are discussed in Bessell (1983) and Bessell <fe Weis cant mass. We need to reassess our knowledge of the known (1987). Transformations between the natural photographic late-M dwarfs to decide how best to follow up these new R/ systems and the Kron-Cousins R/ system are discussed discoveries. by Bessell (1986); the late-M dwarfs are systematically The publication of accurate positions and finding charts of fainter in R ( photographic ) than in R ( photoelectric ), due to faint red and blue high proper-motion stars (the LHS Cata- the broad and asymmetric passband of photoelectric R. log: Luyten 1979; LHS Atlas: Luyten & Albers 1979) stimu- Infrared photometry has also been published by many ob- lated the search for new red and white dwarfs. Reduced mo- servers; Johnson (1965), Veeder (1974), Glass (1974), tions H ( =mR + 5 + 5 log//) were derived for the LHS Mould <fe Hyland ( 1976), Persson etal ( 1977), and Reid & stars of color class m or m +. These were then ordered in Gilmore (1984). Bessell <fe Brett (1987) discuss the rela- decreasing H to provide an observing list. Over the last 10 yr tions between the different JHKL systems and provide trans- we have made spectroscopic and photometric observations, formations to an homogenized system. 662 Astron. J. 101 (2), February 1991 0004-6256/91/020662-15$00.90 © 1991 Am. Astron. Soc. 662 © American Astronomical Society • Provided by the NASA Astrophysics Data System 662B . 663 M. S. BESSELL: THE LATE-M DWARFS 663 101. Standard spectral types obtained from Boeshaar (1976) cussed by Wing (1979a) it is convenient to use the same and Wing (1973, 1979b) have been used for calibration. features for a temperature sequence in both giants and These are identical between Ml and M6.5. Comparisons dwarfs, and in particular because CCD spectra will in the between different system spectral types are discussed by future be even more used, a spectral type based more on the 1991AJ Wing & Yorka (1979). More recently, Turnshek et al. red (7000 to 10 000 Â) TiO and VO bands, CaH and FeH (1985) indicate a dwarf spectral type of MB for VB10 from bands than the yellow CaOH ( dwarf only ) and TiO bands, is digital spectra over the wavelength range 4900-7600 Â. Boe- attractive. This will be discussed in more detail below. shaar ( 1985 ) has extended the dwarf sequence to these later In Table 1 are listed the observed stars, their rough spec- spectral types mainly on the basis of the appearance of the tral types and colors (if available). The FÆ7photometry in bands of CaOH at 5736 Á and VO at 7350 A, and adopted the Cousins system was obtained on the 1 and 2.3 m tele- VB8 and VB10 as the standards for spectral types M7 and scopes at Siding Spring Observatory using the Two-Channel M8, respectively; however, this does produce an abrupt devi- Star-Sky Chopping Photometer and GaAs phototubes. The ation in the relation between giant and dwarf spectral types JHK photometry was obtained on the A AO 3.9 m telescope based on the TiO band strengths as used by Wing. As dis- and the 2.3 m telescope at SSO in the standard A AO and Table 1. Colors, spectral types, reduced motions for LHS stars. LHS GL SP R I K V-R V-I J-H H-K I-K HiR 1002 M5.5 13.77 12.16 10.15 7.42 1.61 3.62 0.60 0.32 2.73 19.5 65B M5.5 12.06 10.40 8.35 5.34 1.66 3.71 0.61 0.35 2.94 19.8 36 406 M6e 13.53 11.67 9.50 6.11 1.86 4.03 0.69 0.36 3.37 21.9 39 412 M5e 14.44 12.77 10.68 7.85 1.67 3.76 0.57 0.27 2.83 22.2 49 551 M5.5e 11.05 9.43 7.43 4.36 1.62 3.62 0.63 0.35 3.06 19.3 68 866 M5.5e 12.37 10.70 8.64 5.66 1.67 3.73 0.66 0.35 2.96 20.4 191 M6.5 18.51 16.24 13.96 10.97 2.21 4.48 0.64 0.35 2.99 22.6 234 283 M6e 16.54 14.68 12.43 9.35 1.86 4.11 0.58 0.32 3.08 21.4 248 1111 M6 14.8 12.8 10.54 7.26 2.00 4.26 0.64 0.37 3.28 20.0 254 M5e 17.2 15.4 13.44 10.78 1.70 3.76 0.51 0.31 2.66 21.9 279 M5.5 20.7 288 M5.5 13.92 12.33 10.31 1.59 3.61 14.9 292 M6 15.73 13.67 11.33 7.97 2.06 4.4 0.66 0.36 3.36 20.9 325 M6e 18.67 16.6 14.36 11.11 2.07 4.31 0.59 0.32 3.25 23.4 330 1159 M6 11.42 0.36 22.6 370 M5 20.9 429 644C M6.5e 16.8: 14.60 12.18 8.85 2.18 4.60 0.64 0.37 3.46 20.9 474 752B M7 17.2: 15.10 12.84 8.81 2.10 4.36 0.67 0.47 4.05 23.7 523 M6e 16.90 14.90 12.56 9.93 2.00 4.34 0.60 0.37 2.63 21.3 1070 M6 15.42 13.71 11.56 8.17 1.71 3.86 0.58 0.43 3.39 18.0 1445 1.45 3.47 20.7 1730 sdM4 21.3 1839 M4.5 21.0 2026 M6e 18.94 16.69 14.32 11.16 2.26 4.61 0.60 0.33 3.16 21.0 2034 M6 22.1 2049 M6 11.41 0.67 0.36 21.2 2065 M7e 18.74 16.74 14.54 9.96 2.00 4.2 0.83 0.51 4.58 21.7 2067 M10* 19.5 15.80 12.84 5.65 3.70 6.60 1.25 0.80 7.19 21.7 2179 M5 21.6 2215 M6 22.2 2236 M5 21.6 2314 M6e 21.6 2347 M5 19.0 17.28 15.12 12.11 1.72 3.87 0.66 0.35 3.01 21.3 2351 M6.5 19.56 17.25 14.91 11.35 2.31 4.65 0.64 0.43 3.58 21.4 2397 M5e 21.3 2397a M7e 19.1 17.1 14.87 10.75 2.00 4.23 0.78 0.45 4.12 21.9 2419 sdM4.: 21.6 2425 M5.5 22.0 2471 M6.5 18.11 16.0 13.66 10.31 2.11 4.45 0.69 0.39 3.35 21.7 2502 M6+ 19.13 17.41 15.33 11.89 1.72 3.80 0.68 0.36 3.44 20.8 2545 M6 22.8 2583 M6 22.7 2589 M4 21.1 2643 M4.5 17.43 16.0 14.15 11.42 1.43 3.28 0.57 0.41 2.73 21.3 2742 M6 22.5 2847 M5 21.9 2855 M6e 19.56 17.6 15.28 11.9 1.96 4.28 0.78 0.47 3.38 22.2 2875 13.04 11.96 10.62 8.55 1.08 2.42 0.63 0.26 2.07 23.7 2876 M6.5 20.14 18.07 15.7 12.09 2.07 4.44 0.64 0.44 3.61 23.7 2924 M7 19.74 15.3 10.68 4.44 0.73 0.49 4.62 22.7 2930 M6.5 17.9 13.3 9.73 4.60 0.66 0.38 3.57 23.8 3001 M5 19.6 3002 M7 22.5 3003 M6.5 17.05 14.88 12.53 8.95 2.17 4.52 0.66 0.42 3.63 21.4 3106 M6 21.5 3189 M5.5 21.2 3307 M6- 21.3 3566 M6 10.58 0.65 0.46 21.2 © American Astronomical Society • Provided by the NASA Astrophysics Data System 662B .