THE SOLAR NEIGHBORHOOD. XXXIII. PARALLAX RESULTS FROM THE CTIOPI 0.9 m PROGRAM: TRIGONOMETRIC PARALLAXES OF NEARBY LOW-MASS ACTIVE AND YOUNG SYSTEMS

All CTIOPI photometry is conducted with the CTIO 0.9 m telescope, initially (1999–2003) under the NOAO Survey Programs grant; later (2003–present) via the SMARTS Consortium. Photometry is conducted in three filters (Tektronix 2 VRI), utilizing only the central quarter (68 × 68 FOV, 401 mas pixel⁻¹) of the Tektronics 2046 × 2046 CCD to minimize distortions for astrometry. These values are then transformed to standard V_JR_KCI_KC¹⁶ (hereafter without subscripts) photometry using observations of standards from Graham (1982), Landolt (1992), and Landolt (2007). The resulting photometry can be found in Table 2. Further details of the observation and reduction procedures can be found in Jao et al. (2005) and Winters et al. (2011). The photometric errors quoted in Columns 3, 4, and 5 of Table 2 combine the Poisson errors, errors on the nightly calibration fit, and standard deviation of multiple nights of photometry (see column 6). Generally, the latter is the greatest contributor to the collective error, particularly when the star is active, as these stars are. This VRI photometry, along with Two Micron All Sky Survey (2MASS) JHK photometry (Skrutskie et al. 2006), is printed in Columns 11, 12, and 13. The measured colors were used to estimate absolute K magnitudes based on the 12 color–magnitude relations presented by Henry et al. (2004). The photometric distances presented in Column 16 of Table 2 were derived from the mean of the distance moduli implied by the absolute K magnitudes and the 2MASS apparent K magnitude.

Relative photometry (for variability studies) comes from our parallax pipeline. With multiple nights of data in the filter used for parallax, we use the methods in Honeycutt (1992) to derive the nightly offsets and zero points for relative instrumental photometry (Jao et al. 2008) to derive stellar variability. These values are given in Column 8 of Table 2.

CTIOPI astrometry is carried out using the same telescope and camera configuration as that used for photometry (Section 3.1) but uses only one filter for each object, chosen to provide the best balance between target(s) and reference star signal-to-noise ratio values. The astrometric pipeline uses all available images taken at hour angles less than two hours, and produces parallaxes, proper motions, and time-series photometry in the parallax filter, all relative to between 5 and 15 "reference" stars within a few arcminutes of the target stars and visible in our images. Parallaxes were corrected to absolute values (Columns 11 and 12 of Table 3) using the mean of the photometric distances to the reference stars, with a typical correction of 1.5 ± 0.5 mas. For a small number of targets with seemingly nearby reference fields (mean photometric parallax estimate, >3.0 mas), we assume the reference stars are actually reddened by some galactic source, and instead apply the typical correction stated above. Between 2005 March and 2009 September, a different V-band filter was used for astrometric and photometric observations. While photometrically identical to the original V filter, it exhibited slightly inferior astrometric performance (Riedel et al. 2010), and all astrometric solutions that incorporate data from it are marked as such in Column 16 of Table 3. Additional details of CTIOPI observing procedures can be found in Jao et al. (2005), Henry et al. (2006), and other papers in this series.

Four of the objects in this paper—BD-21°1074BC, SCR 0613-2742AB, L 449-1AB, and SCR 2010-2801AB—were selected for their X-ray brightness and observed with the Hubble Space Telescope's (HST's) Fine Guidance Sensors (FGSs) in Cycle 16B, in proposal 11943/11944 ("Binaries at the Extremes of the H-R Diagram") using the F583W filter¹⁷ with no pupil. Reductions were carried out for both axes, providing submilliarcsecond precision separations, and delta magnitudes (hereafter Δmag). All four targets were found to be binaries (see Section 5.1 below) and their separations, position angles, and magnitude differences were determined by fitting with single-star fringe scans as described by Nelan et al. (2004).

Hubble's FGSs are implemented as three movable units equipped with Koesters prisms, which allow them to function as a two-dimensional interferometer. Two FGS units lock on guide stars and stabilize the spacecraft, while the third (FGS 1r) scans back and forth across the star. Rather than a Michelson interferometer, light is channeled through a linear polarizing beamsplitter, and then through the unit's Koesters prisms, where the light interferes with itself, allowing two dimensions of fringes to be read out simultaneously.

As an interferometric instrument, FGS is capable of measuring single-axis separations as small as 8 mas, for objects as faint as V = 16.8. It has been routinely used for orbital mapping and submilliarcsecond parallaxes (e.g., Franz et al. 1998; Benedict et al. 1999; McArthur et al. 2011).

3.4.1. CTIO 1.5 m RCSpec

3.4.2. CTIO 4.0 m RCSpec

3.4.3. MPG 2.2 m FEROS

One spectrum of SCR 0613-2742AB was taken with the FEROS spectrograph (Kaufer et al. 1999) on the MPG 2.2 m telescope at La Silla Observatory on 2013 February 18 as part of ESO program 090.C-0200(A). FEROS is an echelle spectrograph fed by two 20 fibers and provides R ∼ 48,000 spectra over a wavelength range of 3500–9200 Å. Observations were taken in the object-sky mode with the use of the atmospheric dispersion corrector. The data were reduced with the facility pipeline and the IRAF task fxcor was used to cross-correlate the target spectrum with several radial velocity standards observed in the same fashion. We measure a heliocentric radial velocity of +22.54 ± 1.16 km s⁻¹ for SCR 0613-2742AB (Table 4). The Na i gravity index measurement and Li 6708 Å EW for SCR 0613-2742AB were also derived from these spectra.

Table 4. Data Used to Calculate UVWXYZ

Name	R.A.	Decl.	π	π	$\mu _{{\rm R.A.\cos decl.}}$	μdecl.	μ	R.V.	R.V.
	(J2000 ± mas)		(mas)	Ref.	(mas yr−1)		Ref.	(km s−1)	Ref.
(1)	(2)	(3)	(4)	(5)	(6)	(7)	(8)	(9)	(10)
NLTT 372	00 08 51.788 ± 70	+20 49 07.86 ± 80	11.66 ± 1.44	1	+120.7 ± 12.1	−76.0 ± 9.4	1	⋅⋅⋅	⋅⋅⋅
G 131-26AB	00 08 53.922 ± 70	+20 50 25.45 ± 80	54.13 ± 1.35	1	−68.5 ± 12.1	−243.0 ± 9.4	2	⋅⋅⋅	⋅⋅⋅
SCR 0017-6645	00 17 23.524 ± 80	−66 45 12.46 ± 110	25.61 ± 1.73	1	+102.9 ± 1.0	−15.0 ± 1.0	2	+11.4 ± 0.8	12
GJ 2006A	00 27 50.242 ± 70	−32 33 06.17 ± 70	30.97 ± 1.76	1 1	+99.2 ± 1.3	−61.3 ± 2.6	2	+8.4	13
GJ 2006B	00 27 50.362 ± 70	−32 33 23.91 ± 70	30.97 ± 1.76	1 1	+117.2 ± 4.1	−31.5 ± 5.8	2	⋅⋅⋅	⋅⋅⋅
SCR 0103-5515ABC	01 03 35.635 ± 60	−55 15 56.19 ± 80	21.18 ± 1.37	1	+100.2 ± 2.0	−47.0 ± 2.4	2	⋅⋅⋅	⋅⋅⋅
LP 467-16AB	01 11 25.412 ± 70	+15 26 21.62 ± 60	45.79 ± 1.78	1	+181.8 ± 12.1	−120.0 ± 9.4	1	+4.0 ± 0.1	14
GJ 2022A	01 24 27.693 ± 60	−33 55 08.75 ± 60	39.50 ± 1.28	1 3	+163.3 ± 12.1	−125.0 ± 9.5	1	+18.4 ± 1.0	3
GJ 2022C	01 24 27.692 ± 60	−33 55 08.75 ± 60	39.50 ± 1.28	1 3	+154.8 ± 12.1	−119.8 ± 9.5	1	+19.4 ± 2.7	3
GJ 2022B	01 24 30.621 ± 60	−33 55 01.64 ± 60	39.50 ± 1.28	1 3	+157.5 ± 12.1	−125.1 ± 9.4	1	+18.3 ± 0.5	3
LP 993-115	02 45 10.711 ± 70	−43 44 32.37 ± 60	87.37 ± 1.33	1 1	+22.7 ± 12.1	−387.2 ± 9.4	1	⋅⋅⋅	⋅⋅⋅
LP 993-116AB	02 45 14.316 ± 70	−43 44 10.60 ± 60	87.37 ± 1.33	1 1	+24.0 ± 12.1	−366.1 ± 9.4	1	⋅⋅⋅	⋅⋅⋅
G 7-34	04 17 18.521 ± 60	+08 49 22.06 ± 60	73.27 ± 1.27	1	+121.6 ± 12.1	−373.8 ± 9.4	1	⋅⋅⋅	⋅⋅⋅
G 39-29AB	04 38 12.592 ± 470	+28 13 00.00 ± 80	75.03 ± 1.89	1 3	+386.9 ± 12.2	−90.7 ± 9.5	1	+35.7 ± 1.2	3
LP 655-48	04 40 23.282 ± 80	−05 30 08.12 ± 70	102.75 ± 0.69	1 3	+330.2 ± 12.1	+127.7 ± 9.4	1	+29.9 ± 0.2	3
LP 476-207ABC	05 01 58.810 ± 60	+09 58 58.77 ± 60	40.18 ± 2.07	1 5	+21.0 ± 12.2	−102.6 ± 9.5	1	+14.7 ± 3.7	15
BD-21°1074BC	05 06 49.472 ± 60	−21 35 03.83 ± 60	51.98 ± 1.30	1 1	+33.1 ± 2.7	−33.2 ± 2.0	2	+23.7 ± 1.7	15
BD-21°1074A	05 06 49.920 ± 60	−21 35 09.19 ± 60	51.98 ± 1.30	1 1	+52.1 ± 1.7	−22.3 ± 1.1	2	+21.2 ± 0.9	15
L 449-1AB	05 17 22.908 ± 60	−35 21 54.70 ± 60	84.38 ± 1.35	1	−234.2 ± 12.1	−162.6 ± 9.5	1	⋅⋅⋅	⋅⋅⋅
SCR 0529-3239	05 29 44.686 ± 60	−32 39 14.17 ± 60	38.19 ± 1.61	1	+12.5 ± 0.9	+13.4 ± 1.8	2	⋅⋅⋅	⋅⋅⋅
SCR 0613-2742AB	06 13 13.308 ± 60	−27 42 05.46 ± 60	34.04 ± 1.00	1	−13.1 ± 1.6	−0.3 ± 1.3	2	+22.54 ± 1.16	1
L 34-26	07 49 12.709 ± 60	−76 42 06.60 ± 60	94.36 ± 2.11	1	−102.4 ± 1.1	−191.9 ± 1.1	2	+0.9	16
SCR 0757-7114	07 57 32.554 ± 60	−71 14 53.81 ± 70	45.25 ± 1.96	1	+86.7 ± 1.3	+20.6 ± 1.3	2	⋅⋅⋅	⋅⋅⋅
SCR 1012-3124AB	10 12 09.085 ± 60	−31 24 45.20 ± 60	18.54 ± 1.74	1	−74.8 ± 1.1	−9.4 ± 1.0	2	+14.69 ± 0.53	13
TWA 8B	11 32 41.165 ± 70	−26 52 09.04 ± 70	21.28 ± 1.01	1 1	−95.3 ± 2.2	−28.6 ± 4.7	2	+8.93 ± 0.27	17
TWA 8A	11 32 41.247 ± 70	−26 51 55.94 ± 70	21.28 ± 1.01	1 1	−95.3 ± 2.2	−28.6 ± 4.7	2	+8.34 ± 0.48	17
SCR 1214-2345	12 14 08.666 ± 80	−23 45 17.06 ± 70	91.39 ± 1.55	1	+42.4 ± 1.1	+85.0 ± 3.4	2	⋅⋅⋅	⋅⋅⋅
G 165-8AB	13 31 46.620 ± 70	+29 16 36.54 ± 60	55.51 ± 2.38	1	−244.1 ± 4.2	−132.4 ± 4.8	2	−7.5 ± 6.5	15
SCR 1425-4113AB	14 25 29.128 ± 60	−41 13 32.40 ± 60	14.94 ± 0.96	1	−46.8 ± 2.1	−49.2 ± 1.7	2	−1.2 ± 1.3	1
GJ 1224	18 07 32.853 ± 60	−15 57 47.05 ± 60	126.99 ± 1.01	1 4	−620.3 ± 12.1	−339.9 ± 9.4	1	−34.8 ± 1.7	15
G 141-29	18 42 44.985 ± 60	+13 54 17.05 ± 70	90.18 ± 1.88	1 4	−39.9 ± 12.1	+359.2 ± 9.5	1	−33.2	15
SCR 1942-2045	19 42 12.818 ± 60	−20 45 47.97 ± 60	62.75 ± 0.90	1 3	−17.2 ± 2.5	−147.9 ± 2.6	2	−1.7 ± 0.2	3
2MASS 2009-0113	20 09 18.242 ± 60	−01 13 38.19 ± 60	95.95 ± 1.54	1	−55.7 ± 12.1	−367.5 ± 9.4	1	⋅⋅⋅	⋅⋅⋅
SCR 2010-2801AB	20 10 00.035 ± 60	−28 01 41.18 ± 60	20.85 ± 1.33	1	+40.7 ± 3.0	−62.0 ± 1.7	2	⋅⋅⋅	⋅⋅⋅
LEHPM2-0783	20 19 49.818 ± 60	−58 16 43.02 ± 60	61.93 ± 1.02	1	−37.2 ± 12.1	−329.7 ± 9.5	1	⋅⋅⋅	⋅⋅⋅
L 755-19	20 28 43.624 ± 60	−11 28 30.82 ± 60	53.18 ± 1.67	1	+155.0 ± 12.1	−91.4 ± 9.5	1	⋅⋅⋅	⋅⋅⋅
SCR 2033-2556	20 33 37.593 ± 60	−25 56 52.15 ± 60	20.70 ± 1.43	1	+52.8 ± 1.7	−75.9 ± 1.3	2	⋅⋅⋅	⋅⋅⋅
SCR 2036-3607	20 36 08.299 ± 60	−36 07 11.48 ± 60	62.13 ± 1.40	1	+5.4 ± 1.7	+43.3 ± 0.9	2	+2	16
GJ 799A	20 41 51.152 ± 130	−32 26 08.00 ± 190	101.31 ± 0.75	1 1 3 3 4 4	+223.2 ± 12.1	−386.9 ± 9.4	1	−3.7 ± 3.0	14
GJ 799B	20 41 51.161 ± 130	−32 26 09.98 ± 190	101.31 ± 0.75	1 1 3 3 4 4	+288.8 ± 12.1	−317.5 ± 9.4	1	−2.7 ± 3.0	14
LHS 3799	22 23 06.997 ± 60	−17 36 26.12 ± 70	137.69 ± 1.75	1 4	+285.9 ± 12.1	−711.3 ± 9.4	1	−2.1 ± 1.1	15
GJ 1284AB	23 30 13.447 ± 60	−20 23 27.44 ± 60	65.79 ± 1.86	1 5	+305.3 ± 12.1	−187.5 ± 9.5	1	−5.7	16
LHS 1302	01 51 04.094 ± 80	−06 07 05.07 ± 80	100.78 ± 1.89	6	+531.8 ± 12.1	−257.9 ± 9.4	6	⋅⋅⋅	⋅⋅⋅
LHS 1358	02 12 54.630 ± 60	+00 00 16.81 ± 70	65.27 ± 2.07	7	+551.2 ± 12.1	+37.8 ± 9.5	7	⋅⋅⋅	⋅⋅⋅
G 99-49	06 00 03.515 ± 60	+02 42 23.63 ± 60	190.77 ± 1.86	4 6	+305.2 ± 12.1	−41.7 ± 9.5	6	+28.7 ± 0.7	15
AP Col	06 04 52.157 ± 60	−34 33 35.98 ± 60	119.21 ± 0.98	8	+21.4 ± 12.1	+341.2 ± 9.4	8	+22.4 ± 0.3	8
G 41-14ABC	08 58 56.332 ± 60	+08 28 26.00 ± 100	147.66 ± 1.98	6	+371.2 ± 3.5	−323.4 ± 2.8	2	−6.4 ± 19.0	15
TWA 27AB	12 07 33.463 ± 60	−39 32 54.01 ± 60	18.95 ± 0.37	9 10 11	−68.5 ± 12.2	−22.5 ± 9.5	9	+7.5 ± 2.0	18
LHS 2729	13 23 38.020 ± 60	−25 54 45.12 ± 70	71.48 ± 1.53	7	−600.9 ± 12.1	−218.3 ± 9.5	7	⋅⋅⋅	⋅⋅⋅
LHS 2836	13 59 10.422 ± 90	−19 50 03.62 ± 60	92.86 ± 0.89	7	−551.2 ± 12.1	−177.0 ± 9.5	7	−15.8	15
GJ 1207	16 57 05.729 ± 60	−04 20 56.29 ± 60	115.26 ± 1.50	4 6	+479.1 ± 12.1	−367.1 ± 9.5	6	−4.2 ± 1.6	15
LHS 4016AB	23 48 36.062 ± 60	−27 39 38.87 ± 70	41.25 ± 1.55	7	−551.9 ± 12.1	−239.8 ± 9.5	7	+25.3 ± 0.55	19

Notes. References: 1: This paper, 2: UCAC4 (Zacharias et al. 2013), 3: Shkolnik et al. (2012), 4: YPC (van Altena et al. 1995), 5: Hipparcos (van Leeuwen 2007), 6: Henry et al. (2006), 7: Riedel et al. (2010), 8: Riedel et al. (2011), 9: Gizis et al. (2007) 10: Biller & Close (2007), 11: Ducourant et al. (2008), 12: Malo et al. (2013), 13: L. Malo et al. (in preparation), 14: Montes et al. (2001), 15: Gizis et al. (2002), 16: Torres et al. (2006), 17: Shkolnik et al. (2011), 18: Rice et al. (2010), 19: Shkolnik et al. (2010). Parallaxes of multiple components in the same system have been combined, and are represented by duplicate references in the parallax column. All positions and position errors are taken from 2MASS (Skrutskie et al. 2006), and adjusted to epoch 2000.0 equinox J2000 using the proper motions listed here.

Download table as:  ASCIITypeset image

Table 5. Deblended Magnitudes for Isochrones

Name	V	I	J	K	MV	V−K	Deblenda
(1)	(2)	(3)	(4)	(5)	(6)	(7)	(8)
NLTT 372	15.87	13.12	11.64	10.72	11.20	5.15
G 131-26A	13.94	10.98	9.40	8.81	12.61	5.13	K
G 131-26B	14.76	11.61	9.90	9.27	13.43	5.49	K
SCR 0017-6645	12.45	10.00	8.56	7.70	9.49	4.75
GJ 2006A	12.95	10.29	8.88	8.01	10.40	4.94
GJ 2006B	13.25	10.48	8.97	8.12	10.70	5.13
SCR 0103-5515A	16.07	12.70	10.83	9.91	12.70	6.16	K
SCR 0103-5515B	16.43	12.97	11.04	10.11	13.06	6.32	K
SCR 0103-5515C	23.39	18.14	14.92	13.69	20.02	9.70	K
LP 467-16A	14.76	11.41	9.53	8.67	13.06	6.09	K
LP 467-16B	15.99	12.36	10.27	9.36	14.29	6.63	K
GJ 2022A	14.27	11.34	9.93	8.97	12.25	5.30	V
GJ 2022C	14.35	11.40	9.98	9.02	12.33	5.33	V
GJ 2022B	15.50	12.33	10.56	9.68	13.48	5.82
LP 993-115	12.38	9.61	8.14	7.27	12.09	5.11
LP 993-116A	13.02	9.58	8.49	7.63	12.73	5.39	i'
LP 993-116B	14.14	10.54	9.27	8.40	13.85	5.74	i'
G 7-34	13.84	10.75	9.03	8.18	13.16	5.66
G 39-29A	13.03	10.23	8.75	7.91	12.41	5.12	K
G 39-29B	13.69	10.74	9.14	8.28	13.07	5.41	K
LP 655-48	17.79	13.36	10.66	9.55	17.85	8.24
LP 476-207A	12.46	9.75	8.30	7.48	10.48	4.98	SB, K
LP 476-207B	13.56	10.41	8.65	7.76	11.58	5.80	K
LP 476-207C	12.46	9.75	8.30	7.48	10.48	4.98	SB, K
BD-21°1074B	11.44	8.86	7.47	6.60	10.02	4.84	R
BD-21°1074C	12.44	9.71	8.14	7.21	11.02	5.23	R
BD-21°1074A	10.41	8.25	7.05	6.12	8.99	4.29
L 449-1A	12.05	9.32	7.87	7.05	11.68	5.00	R
L 449-1B	13.06	10.17	8.54	7.66	12.69	5.40	R
SCR 0529-3239	13.79	10.80	9.22	8.32	11.70	5.47
SCR 0613-2742A	12.78	10.07	8.57	7.72	10.44	5.06	R
SCR 0613-2742B	13.41	10.60	8.98	8.11	11.07	5.30	R
L 34-26	11.31	8.79	7.41	6.58	11.18	4.73
SCR 0757-7114	12.45	9.77	8.32	7.42	10.73	5.03
SCR 1012-3124A	14.25	11.25	9.59	8.74	10.59	5.51	V
SCR 1012-3124B	14.28	11.27	9.61	8.75	10.62	5.53	V
TWA 8B	15.22	11.76	9.84	9.01	11.86	6.21
TWA 8A	12.23	9.79	8.34	7.43	8.87	4.80
SCR 1214-2345	13.96	10.78	9.07	8.23	13.80	5.73
G 165-8A	12.64	9.80	8.23	7.40	11.36	5.24	K
G 165-8B	12.93	10.02	8.40	7.56	11.65	5.37	K
SCR 1425-4113A	13.29	10.81	9.30	8.37	9.16	4.92	SB
SCR 1425-4113B	13.29	10.81	9.30	8.37	9.16	4.92	SB
GJ 1224	13.48	10.31	8.64	7.83	14.00	5.65
G 141-29	12.86	9.95	8.36	7.55	12.64	5.31
SCR 1942-2045	14.33	11.25	9.60	8.76	13.31	5.57
2MASS 2009-0113	14.47	11.16	9.40	8.51	14.38	5.96
SCR 2010-2801A	13.62	10.85	9.32	8.41	10.22	5.21	R
SCR 2010-2801B	13.86	11.06	9.49	8.56	10.46	5.30	R
LEHPM2-0783	17.17	13.03	10.66	9.72	16.13	7.46
L 755-19	12.47	9.81	8.39	7.50	11.10	4.97
SCR 2033-2556	14.87	11.57	9.71	8.88	11.45	5.99
SCR 2036-3607	11.66	9.27	8.03	7.17	10.63	4.49
GJ 799A	11.09	8.14	6.55	5.69	11.12	5.40	V
GJ 799B	11.13	8.17	6.57	5.71	11.16	5.42	V
LHS 3799	13.30	10.04	8.24	7.32	13.99	5.98
GJ 1284A	11.89	9.34	7.95	7.08	10.98	4.81	SB
GJ 1284B	11.89	9.34	7.95	7.08	10.98	4.81	SB
LHS 1302	14.49	11.16	9.41	8.55	14.51	5.94
LHS 1358	13.58	10.66	9.06	8.17	12.65	5.41
G 99-49	11.31	8.43	6.91	6.04	12.71	5.27
AP Col	12.96	9.60	7.74	6.87	13.34	6.09
G 41-14A	12.07	9.27	7.78	6.97	12.92	5.10	K, SB
G 41-14B	12.07	9.27	7.78	6.97	12.92	5.10	K, SB
G 41-14C	12.21	9.20	7.56	6.72	13.06	5.49	K
TWA 27A	19.95	15.92	13.00	11.96	16.34	7.99	K
TWA 27B	28.91	22.74	18.33	16.94	25.30	11.97	K
LHS 2729	12.89	10.14	8.66	7.78	12.16	5.11
LHS 2836	12.88	9.90	8.33	7.45	12.72	5.44
GJ 1207	12.25	9.43	7.97	7.12	12.56	5.13
LHS 4016A	13.09	10.65	9.34	8.50	11.17	4.59	SB
LHS 4016B	13.09	10.65	9.34	8.50	11.17	4.59	SB

Note. ^aBand of known Δmag used to deblend the photometry, or SB if system was identified only as a spectroscopic binary. No entry is given for single stars because deblending is not required.

Download table as:  ASCIITypeset images: 1 2

3.4.4. CFHT ESPaDOnS

Stars that form together should be moving together through space. Over time, internal and external interactions will cause them to disperse into the thin disk of the Galaxy. The associations we consider here (Table 8) are all sufficiently young that this has not happened yet, even though (with the exception of η Chameleontis, the Pleiades, and the Hyades) they are gravitationally unbound.

Table 8. Nearby Young Associations

Name	U	σU	V	σV	W	σW	min X	max X	min Y	max Y	min Z	max Z	Age
	(km s−1)	(km s−1)	(km s−1)	(km s−1)	(km s−1)	(km s−1)	(pc)	(pc)	(pc)	(pc)	(pc)	(pc)	(Myr)
(1)	(2)	(3)	(4)	(5)	(6)	(7)	(8)	(9)	(10)	(11)	(12)	(13)	(14)
Cha	−11.0	1.2	−19.9	1.2	−10.4	1.6	+34	+60	−105	−78	−44	−12	6a
TW Hya	−10.5	0.9	−18.0	1.5	−4.9	0.9	+2	+34	−61	−26	+10	+27	8a
β Pic	−10.1	2.1	−15.9	0.8	−9.2	1.0	−32	+76	−33	21	−29	−1	12a
Octans	−14.5	0.9	−3.6	1.6	−11.2	1.4	−79	+142	−138	−60	−85	−38	20
Tuc-Hor	−9.9	1.5	−20.9	0.8	−1.4	0.9	−61	+43	−47	−4	−44	−30	30
Columba	−13.2	1.3	−21.8	0.8	−5.9	1.2	−106	+9	−168	+1	−99	+6	30
Carina	−10.2	0.4	−23.0	0.8	−4.4	1.5	−2	+33	−154	−39	−33	+5	30
Argus	−22.0	0.3	−14.4	1.3	−5.0	1.3	−55	+64	−154	−6	−67	+8	50
AB Dor	−6.8	1.3	−27.2	1.2	−13.3	1.6	−94	+73	−131	+58	−66	+23	125b
Pleiadesc	−6.6	0.4	−27.6	0.3	−14.5	0.3	−134	−108	+14	+40	−66	−40	125b
Castord	−10.7	3.5	−8.0	2.4	−9.7	3.0	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	200
UMae	+14.56	2.28	+2.81	1.75	−8.37	3.42	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	500
Hyadesf	−41.1	2.0	−19.2	2.0	−1.4	2.0	−53	−33	−9.3	+10.7	−27.3	+7.3	650

Notes. All data (unless otherwise specified) from Torres et al. (2008). ^aOnly these ages are known with any degree of certainty or corroboration. ^bAge from Luhman et al. (2005). ^cDimensions Soderblom et al. (2005), 13 pc tidal radius Adams et al. (2001). ^dBarrado y Navascues (1998). ^eKing et al. (2003). ^fRöser et al. (2011).

Download table as:  ASCIITypeset image

The perhaps hundreds of systems in an association are spread out over tens of thousands of cubic parsecs, interspersed among thousands of field systems (and, indeed, members of other young moving groups), and as an unbound association, their velocity dispersions are larger than 1 km s⁻¹. Therefore, large numbers of older stars (López-Santiago et al. 2009 in particular quotes ∼30%) will coincidentally happen to have matching UVW motions, and even larger numbers of field stars will fall within the spatial boundaries of an association. UVW motions do not prove youth, but they are necessary to connect young stars to young associations.

There is also the non-trivial issue of whether we have an accurate assemblage of these nearby associations and moving groups. We adopt the associations in Table 8 despite the knowledge that the physical reality of these groups (and the accuracy of their proposed members) is still somewhat uncertain. For instance, the IC 2391 Supercluster (Eggen 1991), Carina-Vela moving group (Makarov & Urban 2000), and Argus association (Torres et al. 2008) have all been proposed as the extended halo of the nearby IC 2391 open cluster, but all are more or less distinct from each other in terms of membership and proposed properties; currently only Argus is thought to be an actual co-eval assembly.

4.1.1. The Kinematic Data

4.1.2. The Kinematic Method

The standard method for computing UVW space velocities is laid out in Johnson & Soderblom (1987), and the matrices in that paper can easily be adapted to compute XYZ space positions. These UVWXYZ coordinates are right-handed Cartesian Galactic coordinates aligned so that the U/X axis is toward the galactic center, the V/Y axis is in the direction of galactic rotation, and the W/Z axis is toward the north Galactic pole. In cases where we have full kinematic information, we calculate 10⁷ Monte Carlo iterations to fully sample the uncertainties on our input kinematics as a three-dimensional ellipsoid in velocity space. For stars without radial velocity measurements, we calculate 10⁵ Monte Carlo iterations at multiple different radial velocities within a range −100 to +100 km s⁻¹ (see Figure 2). UVWXYZ coordinates for the target objects are given in Table 9.

Table 9. UVWXYZ Kinematics

Name	U	V	W	X	Y	Z
	(km s−1)	(km s−1)	(km s−1)	(pc)	(pc)	(pc)
(1)	(2)	(3)	(4)	(5)	(6)	(7)
NLTT 372	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	−21.86 ± 2.73	+61.00 ± 7.61	−56.25 ± 7.01
G 131-26AB	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	−4.71 ± 0.12	+13.13 ± 0.33	−12.10 ± 0.30
SCR 0017-6645	−11.30 ± 1.12	−16.71 ± 0.87	−9.19 ± 0.63	+15.49 ± 1.05	−19.72 ± 1.33	−29.92 ± 2.02
GJ 2006A	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+4.04 ± 0.23	−1.10 ± 0.06	−32.01 ± 1.80
GJ 2006B	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+4.05 ± 0.23	−1.10 ± 0.06	−32.02 ± 1.82
SCR 0103-5515ABC	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+10.93 ± 0.70	−19.50 ± 1.26	−41.64 ± 2.68
LP 467-16AB	−12.34 ± 1.16	−16.58 ± 1.19	−9.87 ± 0.72	−9.55 ± 0.37	+11.37 ± 0.44	−16.01 ± 0.62
GJ 2022A	−7.71 ± 1.36	−26.52 ± 1.47	−13.93 ± 1.02	−0.85 ± 0.03	−4.24 ± 0.14	−24.95 ± 0.81
GJ 2022C	−6.72 ± 1.36	−25.58 ± 1.51	−15.12 ± 2.67	−0.85 ± 0.03	−4.24 ± 0.14	−24.95 ± 0.81
GJ 2022B	−6.55 ± 1.36	−26.10 ± 1.45	−13.92 ± 0.56	−0.86 ± 0.03	−4.24 ± 0.14	−24.95 ± 0.81
LP 993-115	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	−1.28 ± 0.02	−5.27 ± 0.08	−10.08 ± 0.16
LP 993-116AB	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	−1.28 ± 0.02	−5.27 ± 0.08	−10.08 ± 0.15
G 7-34	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	−11.95 ± 0.21	−0.96 ± 0.02	−6.52 ± 0.11
G 39-29A	−40.54 ± 1.19	−14.62 ± 0.84	+6.88 ± 0.82	−12.89 ± 0.33	+1.84 ± 0.05	−2.87 ± 0.07
G 39-29B	−32.71 ± 1.57	−15.74 ± 0.86	+8.62 ± 0.85	−12.89 ± 0.33	+1.84 ± 0.05	−2.87 ± 0.07
LP 655-48	−30.67 ± 0.33	−14.73 ± 0.47	−1.84 ± 0.47	−7.69 ± 0.05	−3.12 ± 0.02	−5.09 ± 0.03
LP 476-207ABC	−10.12 ± 3.47	−13.55 ± 1.50	−9.10 ± 1.77	−23.15 ± 1.19	−4.30 ± 0.22	−8.06 ± 0.42
BD-21°1074BC	−12.92 ± 1.07	−17.06 ± 0.99	−11.06 ± 0.93	−12.01 ± 0.30	−11.00 ± 0.28	−10.25 ± 0.26
BD-21°1074A	−12.37 ± 0.57	−16.07 ± 0.54	−8.06 ± 0.50	−12.01 ± 0.30	−11.00 ± 0.28	−10.24 ± 0.26
L 449-1AB	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	−5.07 ± 0.08	−8.47 ± 0.14	−6.56 ± 0.11
SCR 0529-3239	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	−12.41 ± 0.52	−18.83 ± 0.79	−13.28 ± 0.56
SCR 0613-2742AB	−12.43 ± 0.65	−16.42 ± 0.90	−9.34 ± 0.45	−15.99 ± 0.47	−22.49 ± 0.66	−10.07 ± 0.30
L 34-26	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+3.16 ± 0.07	−9.22 ± 0.21	−4.16 ± 0.09
SCR 0757-7114	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+4.91 ± 0.21	−20.11 ± 0.87	−7.74 ± 0.34
SCR 1012-3124A	−14.59 ± 1.34	−17.80 ± 0.63	−7.61 ± 1.23	−2.38 ± 0.22	−50.55 ± 4.75	+18.65 ± 1.75
SCR 1012-3124B	−14.57 ± 1.34	−17.55 ± 0.80	−7.70 ± 1.24	−2.38 ± 0.22	−50.55 ± 4.75	+18.65 ± 1.75
TWA 8B	−14.10 ± 0.98	−18.15 ± 0.74	−6.54 ± 1.01	+8.21 ± 0.39	−38.65 ± 1.84	+25.44 ± 1.21
TWA 8A	−14.20 ± 0.39	−17.67 ± 1.83	−6.85 ± 1.03	+8.21 ± 0.39	−38.65 ± 1.83	+25.44 ± 1.21
SCR 1214-2345	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+3.22 ± 0.06	−7.96 ± 0.14	+6.79 ± 0.12
G 165-8AB	−10.17 ± 0.87	−22.38 ± 1.26	−3.81 ± 6.42	+1.88 ± 0.08	+2.16 ± 0.09	+17.79 ± 0.76
SCR 1425-4113AB	−10.95 ± 1.24	−16.35 ± 1.43	−8.85 ± 0.86	+49.64 ± 3.19	−39.70 ± 2.55	+20.97 ± 1.35
GJ 1224	−28.98 ± 1.65	−30.09 ± 0.57	+12.77 ± 0.45	+7.64 ± 0.06	+1.87 ± 0.02	+0.30 ± 0.01
G 141-29	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+7.84 ± 0.16	+7.68 ± 0.16	+1.58 ± 0.03
SCR 1942-2045	+1.08 ± 0.20	−11.11 ± 0.25	−2.28 ± 0.20	+14.14 ± 0.20	+4.94 ± 0.07	−5.50 ± 0.08
2MASS 2009-0113	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+7.50 ± 0.13	+6.51 ± 0.11	−3.20 ± 0.05
SCR 2010-2801AB	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+40.83 ± 2.63	+10.28 ± 0.66	−22.93 ± 1.48
LEHPM2-0783	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+12.46 ± 0.21	−4.76 ± 0.08	−9.11 ± 0.15
L 755-19	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+14.02 ± 0.44	+9.24 ± 0.29	−8.46 ± 0.27
SCR 2033-2556	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+38.49 ± 2.66	+12.69 ± 0.88	−26.32 ± 1.82
SCR 2036-3607	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+12.96 ± 0.29	+1.45 ± 0.03	−9.43 ± 0.21
GJ 799A	−7.99 ± 2.40	−17.47 ± 0.65	−9.01 ± 1.84	+7.80 ± 0.06	+1.54 ± 0.01	−5.85 ± 0.04
GJ 799B	−9.30 ± 2.40	−13.86 ± 0.65	−11.49 ± 1.84	+7.80 ± 0.06	+1.54 ± 0.01	−5.85 ± 0.04
LHS 3799	−0.02 ± 0.61	−24.92 ± 0.60	−8.95 ± 0.94	+3.20 ± 0.04	+2.74 ± 0.04	−5.92 ± 0.08
GJ 1284AB	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+3.40 ± 0.10	+3.85 ± 0.11	−14.31 ± 0.40
LHS 1302	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	−3.98 ± 0.08	+1.48 ± 0.03	−8.97 ± 0.17
LHS 1358	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	−8.04 ± 0.25	+2.60 ± 0.08	−12.78 ± 0.40
G 99-49	−24.82 ± 0.64	−16.29 ± 0.37	+1.05 ± 0.31	−4.69 ± 0.05	−2.15 ± 0.02	−0.92 ± 0.01
AP Col	−22.00 ± 0.37	−13.46 ± 0.36	−4.65 ± 0.45	−3.72 ± 0.03	−6.70 ± 0.06	−3.41 ± 0.03
G 41-14ABC	+16.03 ± 12.28	−5.62 ± 10.37	+1.62 ± 10.10	−4.38 ± 0.06	−3.70 ± 0.05	+3.60 ± 0.05
TWA 27A	−10.33 ± 2.88	−15.76 ± 2.30	−5.14 ± 2.35	+19.66 ± 0.38	−44.60 ± 0.87	+20.23 ± 0.40
LHS 2729	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+7.53 ± 0.16	−8.38 ± 0.18	+8.30 ± 0.18
LHS 2836	⋅⋅⋅	⋅⋅⋅	⋅⋅⋅	+6.65 ± 0.06	−4.83 ± 0.05	+6.96 ± 0.07
GJ 1207	+6.19 ± 1.44	−0.83 ± 0.57	−24.39 ± 0.82	+7.72 ± 0.10	+2.05 ± 0.03	+3.38 ± 0.04
LHS 4016AB	+72.89 ± 2.85	+4.57 ± 1.15	−9.44 ± 0.85	+5.20 ± 0.20	+2.66 ± 0.10	−23.53 ± 0.88

Notes. Six-dimensional phase space positions for the stars in this sample.

Download table as:  ASCIITypeset image

To determine whether a star system is a potential match for a given association, we must determine how close the three-dimensional velocity–space ellipsoid(s) defined by our Monte Carlo iterations is to the three-dimensional velocity–space ellipsoid of the association, as defined in Table 8. Because the dispersions are meaningful—in the case of the sample system, they represent the uncertainty on the velocity of the system; in the case of the association, they represent the intrinsic dispersion of real members—our phase-space "separations" are calculated relative to those dispersions, in the form of our goodness-of-fit statistic γ:

$$\begin{eqnarray} \gamma &=& \frac{1}{3} \left(\frac{(U_{{\rm assoc}}-U_{{\rm system}})^{2}}{\left(\sigma _{U_{{\rm assoc}}}^{2} + \sigma _{U_{{\rm system}}}^{2}\right)} + \frac{(V_{{\rm assoc}}-V_{{\rm system}})^{2}}{\left(\sigma _{V_{{\rm assoc}}}^{2} + \sigma _{V_{{\rm system}}}^{2} \right)}\right. \nonumber\\ &&\left. + \frac{(W_{{\rm assoc}}-W_{{\rm system}})^{2}}{\left(\sigma _{W_{{\rm assoc}}}^{2} + \sigma _{W_{{\rm system}}}^{2}\right)} \right). \end{eqnarray} \tag{ 1 }$$

This γ statistic is effectively identical to the one used in Shkolnik et al. (2012), where it appears as $\bar{\chi }^2$.

We take a value of γ less than 4 as a potentially significant match. In many cases, a system will be kinematically consistent with membership in more than one young association. This is often unavoidable, as the velocity distributions of several associations genuinely overlap; in these cases we must look at the other diagnostics to determine which association is the most consistent with the available data. In cases where we have no radial velocity, we fit the resulting γ values from our range of points to determine a best-fit radial velocity and γ. Only the γ value for the most consistent association is given in Column 5 of Table 7.

Pre-main-sequence stars are still in the process of contracting under gravity, and are still physically larger than main-sequence stars of the same temperature, and consequently much brighter. It is therefore possible to distinguish stars of a given age using isochrones, or at least demonstrate the youth of a system.

4.2.1. The Photometric Data

Overluminosity is also a sign of multiplicity, and must therefore be taken into account before making conclusions about the potential youth of a system based on its luminosity. Our sample has many multiple stars (see Table 10), which are not resolved in CTIOPI images if they are less than ∼2'' apart. In order to properly distinguish between binaries (which in the maximal case of an equal-luminosity pair can be 0.7 mag brighter than a single star) and young stars, we have made an extensive literature search for multiples within the sample. Once multiples are identified, their photometry must be deblended to properly place them on a color–magnitude diagram.

Unfortunately, we do not have both ΔV and ΔK_S values for any multiples in our sample, so the deblending in one of the filters must be estimated. Plotting, for instance, M_V versus $M_{K_{S}}$ (Figure 3) demonstrates that the relation is, to a first approximation, linear between M_V = 1 and M_V = 15 ($M_{K_S}$ = 1 and $M_{K_S}$ = 9), in that case with a slope of 1.8, for ΔV = 1.8 × ΔK_S. We list all the equivalencies in Table 11. Actual measurements should definitely be preferred, however, as the error on the fit slopes are on the order of 0.03, and the residuals to the fits are on the order of 0.5.

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

For purposes of deblending our stars, we have assumed that the ΔFGS583W values are equivalent to Kron–Cousins ΔR, and that all the various ΔK values (CIT, Altair, MKO, 2MASS) are equivalent to 2MASS ΔK_S. Several systems are known only as spectroscopic binaries. With no other information available, we have assumed ΔV = ΔK_S = 0. The photometric data used in the isochrone analysis, including deblended magnitudes, are given in Table 5 and used to place stars on the color–magnitude diagram in Figure 4. Deblended optical magnitudes are given for the brown dwarfs TWA 27B and SCR 0103-5515C, but their deblended V magnitudes are suspect at best.

Zoom In Zoom Out Reset image size

4.2.2. The Photometric Method

BANYAN (Malo et al. 2013) is an independent Bayesian methodology for finding young stars. BANYAN uses IJ photometry, Baraffe et al. (1998, 2002) model isochrones, and a slightly different set of UVWXYZ values for the known associations. It searches for members of the known nearby associations β Pictoris, TW Hydra, Tucana-Horologium, Columba, Carina, Argus, and AB Doradus, with "Field" as the default hypothesis. It is presented as a complementary, proven method. The coefficients for the best-matching association are shown in Table 7, Columns 5 and 6.

Three gravity-sensitive features exist within the 6000 Å–9000 Å coverage of our spectra: Ca ii (8498,8542,8662 Å) is strong in giants and weak in dwarfs; Na i (8183,8195 Å) and K i (7665,7699 Å) are weak in giants and strong in dwarfs. The general pattern is that the neutral alkali and alkali earth metals are increasingly strong with higher gravity, while their singly ionized variants grow weaker with higher gravity (Allers et al. 2007; Schlieder et al. 2012a).

In this paper we are using the Na i index as defined in Lyo et al. (2004a), and the K i 7699 Å doublet line (the other is contaminated by the atmospheric A band). The Na i index value is formed by the ratio of the average flux in two 24 Å wide bands:

$$\begin{equation} {\rm Na\,\scriptsize{I}}_{{\rm index}} = \frac{F_{8148\hbox{--}8172}}{F_{8176\hbox{--}8200}}. \end{equation} \tag{ 2 }$$

Measurements for the program stars are given in Table 6, and graphs of the general trends are shown in Figure 5. Unfortunately, as can be seen in the figures, young stars and main-sequence stars overlap at colors bluer than V − K_s = 5, in both cases. The lines can also be affected by stellar activity, where emission fills in the absorption line cores, leading to lower EWs (Reid & Hawley 1999).

Zoom In Zoom Out Reset image size

In principle, multiplicity and metallicity will have an effect on these features. The effects are muted in the case of multiplicity, as the brighter component will dominate, and the line will not be appreciably weaker or stronger than that of a main-sequence star of the same color. Metallicity is more difficult. While we can expect low-metallicity subdwarfs to have weaker lines due to lower abundances, subdwarfs plot below the main sequence and will not be mistaken for pre-main-sequence stars. However, as noted by Shkolnik et al. (2009), high-metallicity stars of a given mass and bolometric luminosity will masquerade as stars of lower temperature and lower surface gravity. The additional metals increase the opacity of the stellar atmosphere and therefore put the effective photosphere farther from the center of the star. Thus, these stars will appear above the main sequence—in fact, these gravity-sensitive atomic species are the ones used by Rojas-Ayala et al. (2010) to measure the metallicity of field M dwarfs.

With the exception of SCR 0613-2742AB and SCR 1425-4113AB, the only available spectra for the sample stars are low-resolution flux-calibrated optical spectra from the CTIO 1.5 m telescope that cannot be reliably corrected for telluric absorption. Thus, there is an additional source of error in our Na i index measurements, and we are using only one of the K i doublet lines (the other is within the atmospheric A band). We conclude that while we see some indication of offsets for young stars in the plots of Figure 5, conclusions of youth using this method with data at this resolution are only tentative.

4.5.1. X-Ray Activity

4.5.2. Hα Emission

M dwarfs also have saturated Hα emission for long periods of time (West et al. 2008). We have measured this value (Table 6), but it cannot be used to distinguish between the ages of low-mass pre-main-sequence stars. The only useful purpose of Hα emission for low-mass stars is to provide an upper (if present) or lower (if absent) limit on ages, which has been done in Table 7 following the prescription in West et al. (2008). Though we see substantial Hα emission in a few of our stars, none of their emissions reach the White & Basri (2003) threshold necessary to be considered a T Tauri star (see Figure 6).

Zoom In Zoom Out Reset image size

4.5.3. Photometric Variability and Flares

One of the hallmarks of the T Tauri class of young stars is variability, and this extends into the older non-accreting stars discussed here. Analysis of the relative variability of M dwarfs (Jao et al. 2011) shows that a typical M dwarf varies by roughly 0.010 mag, and statistically significantly more in V and R (0.013 mag) than I (0.008 mag). They also found a statistically significant difference between regular M dwarfs and their older, metal poor subdwarf cousins (variability 0.007 mag, our observational "floor") which points to some combination of age and (possibly) metallicity influencing the amplitude of stellar variability. As seen in Figure 7, M dwarfs with variability higher than 0.020 mag are rare, although many are present among the young stars in this paper.

Zoom In Zoom Out Reset image size

While many stars in this paper are known flare stars or UV Ceti variables, only two stars were seen to flare during the course of astrometric observations: TWA 8B and GJ 1207. The relative photometry during the flares is reproduced in Figure 8.

Zoom In Zoom Out Reset image size

NLTT 372 was originally part of the reference field of the system G 131-26AB. It is within the error circle of the ROSAT X-ray source we attribute to G 131-26AB, and may either contribute or be the source of those X-rays. It is more luminous than a single star of its photometric colors, and is likely to be a binary.

GJ 2006AB appear to be β Pictoris members, based on kinematics, gravity, and their positions relative to the β Pictoris isochrone. A spectrum of GJ 2006A was obtained with Very Large Telescope (VLT) instrumentation (L. Malo et al., in preparation) and the radial velocity was determined to be +8.90 km s⁻¹ with vsin i of 6 km s⁻¹, in excellent agreement with that of a predicted β Pictoris member.

SCR 0103-5515ABC was resolved as a close triple by Delorme et al. (2013) composed of two M dwarfs, A and B, and a more spatially distant brown dwarf, C (see Table 10). Delorme et al. (2013) and the BANYAN results find it to be a match to Tucana-Horologium, but we find that its kinematics better match the Carina association, and because both associations are supposed to be the same age, we cannot distinguish between them with the other methods. When their photometry is deblended, both components lie above the β Pictoris isochrone, implying that they are younger than β Pic, even though Tucana-Horologium and Carina are both supposed to be older than β Pic. The system is undoubtedly young, but more observations are needed to determine if it is a higher-order multiple in Tuc-Hor, or something entirely different.

GJ 2022ABC is a hierarchical triple (see Table 10) composed of a wide (378) companion (B) to a close (18) nearly equal-luminosity pair (AC) with a delta magnitude of roughly ΔV = 0.08 (Jao et al. 2003; Daemgen et al. 2007). The B component is actually the least luminous, though to preserve the historical order of discovery, we continue to refer to it as "B." The AC component is bright in RASS; there is no corresponding detection of the B component.

The A and C components are separated by 18 (Figure 9). They are separable and (as they are much brighter than the B component) unsaturated on only 31 of the 66 images taken. Those 31 images were taken on 14 nights spanning 11.89 yr and positions were obtained using special extra-sensitive (but less precise) SExtractor settings. A reduction of all three components yields

• 1.  
  GJ 2022A: π = 40.91 ± 5.64 mas, μ = 210.5 ± 1.3 mas yr⁻¹ @ 1265 ± 071
• 2.  
  GJ 2022B: π = 42.12 ± 3.60 mas, μ = 206.2 ± 0.8 mas yr⁻¹ @ 1274 ± 046
• 3.  
  GJ 2022C: π = 46.96 ± 5.31 mas, μ = 200.6 ± 1.2 mas yr⁻¹ @ 1268 ± 071.

Zoom In Zoom Out Reset image size

This implies a weighted mean parallax of 43.05 ± 2.63 mas (23.2 ± 1.4 pc), which is consistent with our main reduction of B in Table 3 using 66 frames (38.80 ± 2.13 mas, 25.8 ± 1.4 pc) and with membership in the AB Dor association. Formally, we are using the parallax of the B component alone as our system parallax in Table 4 due to the lower precision of the results for all three components.

LHS 1302 is potentially a kinematic match for the ∼30 Myr old Columba association, but its gravity (Figure 5) and isochrones (Figure 4) suggest it is a field object.

LP 993-115 (A)/LP 993-116AB (BC). We detect the astrometric orbital motion of the BC pair (Figure 10). Our attempt to fit an orbit did not converge, and the astrometry in Table 3 was calculated without compensating for orbital motion. This triple system appears to be composed of field stars, and only the BC components have X-ray emission.

Zoom In Zoom Out Reset image size

LP 476-207ABC was observed by Hipparcos as HIP 23418; the resulting parallax was of poor quality due to erroneous coordinates in the input catalog (Perryman et al. 1997) and the resulting poor available astrometry. Our new result (24.6 ± 1.3 pc) is significantly closer and higher precision than the old (Perryman et al. 1997; 32.1 ± 8.8 pc) and revised (van Leeuwen 2007; 33.2 ± 10.5 pc) Hipparcos values, but the system is still in β Pic.

BD-21°1074ABC is a known member of the β Pic association (Torres et al. 2008; Malo et al. 2013). The A component was not originally targeted for parallax measurement and was saturated in many early images; the parallax result in Table 3 is of lower precision.

The BC component was observed by the HST FGS Interferometer on 2008 December 18, and resolved (Figure 11) in both axes, with details in Table 10. The measured position angle is discrepant with the current Washington Double Star catalog value (08 @ 321°) but matches the angle of the visible elongation of the BC point spread function from CTIOPI data, as seen in Figure 12. We see an astrometric binary signal in our astrometry, as shown in Figure 13.

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

L 449-1AB was identified as an active star by Scholz et al. (2005). This system was on the HST Cycle 16B FGS proposal, and was resolved into two stars on 2008 December 3 (Figure 14) with ΔF583W = 0.95 mag (see Table 10). The astronomic data displays what may be a very weak perturbation (Figure 15) due to orbital motion, but it is not convincing. By H-R diagram isochrones and gravity indices, this system is indistinguishable from main-sequence stars. The system's kinematics are a potential match for the Ursa Major moving group, which given the probable age of 500 Myr (King et al. 2003) again suggests that the components may be nearly zero-age main sequence (ZAMS).

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

LHS 1358 has kinematics consistent with the Hyades stream. Various authorities (e.g., Famaey et al. 2008) dispute the existence of the stream as a real kinematic entity; at only 15.3 ± 0.5 pc from the Sun, LHS 1358 is a minimum of 30 pc (3 tidal radii) from the Hyades cluster center, and therefore the identification is not physical.

SCR 0613-2742AB is the lowest proper motion system (11.2 ± 1.0 mas yr⁻¹) thus far observed on CTIOPI, with a transverse velocity of 1.6 km s⁻¹. The system is a binary, observed on 2008 December 4 and resolved into two stars (Figure 16, Table 10) by FGS, and is also detected as a binary by our FEROS spectroscopy. We see orbital motion in the CTIOPI astrometric residuals of this star (Figure 17) and have removed that orbital motion¹⁹ from our astrometric results (Riedel et al. 2010), but we do not have data on a full orbit. Based on the FEROS spectroscopy, we see no Li 6708 Å doublet absorption. This is expected for β Pictoris members near the lithium depletion boundary (Yee & Jensen 2010).

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

The position and proper motion of SCR 0613-2742AB, near the convergent point of the β Pic association, are extremely similar to those of 2MASS J06085283–2753583 (Cruz et al. 2003; Rice et al. 2010), a brown dwarf β Pic member. The parallax of 2M0608 (32.0 ± 3.6 mas, 31.3 ± 3.5 pc; Faherty et al. 2012) is also very similar to SCR0613 (34.0 ± 1.0 mas, 29 ± 0.9 pc). The two systems are separated by 3529'' @ 786, which yields a minimum projected separation of 100 kAU (≈0.5 pc). It is difficult to compare their proper motions—(+8.9,+10.7) ± (3.5) mas yr⁻¹ relative (2M0608) and (−13.1,−0.3) ± (11.6,15.5) mas yr⁻¹ absolute (SCR0613)—given the large uncertainties involved, save that each are marginally consistent with both zero and each other.

Meanwhile, the binding energy, U_g = (− GM₁M₂/r) = −1.28 × 10³³J using the conservative estimates that the two components of SCR0613 have a combined mass of 0.5 M_☉, and 2M0608 is 0.015 M_☉. This binding energy is lower than any of the extremely wide systems studied in Caballero (2009), and it is probable that these stars are not gravitationally bound. It is possible that this was a bound system at one point (like AU/AT Mic = GJ 803/799AB, below), but is not any longer.

SCR 0757-7714 is overluminous by 2 mag, plotting near the β Pic isochrone on the H-R diagram in Figure 4, and has low surface gravity based on its K i measurement, but not Na i. Surprisingly, it shows no other signs of youth or membership in any known association—in fact, it has no X-ray emission, and is the only star under consideration with Hα in absorption. Thus, the star's elevation on the H-R diagram is most likely due to unresolved multiplicity.

L 34-26 is a potential kinematic match to the Ursa Major moving group, but its measured radial velocity (+0.9 km s⁻¹, no error; Torres et al. 2006) is most likely discrepant with the best-fit Ursa Major moving group radial velocity, +6.6 km s⁻¹.

SCR 1012-3124AB is a close ∼1'' binary (Figure 18 and Table 10). There are a few epochs where the B component (to the west of the A component) can be seen, but for the most part the two components are blended in our astrometric observations.

Zoom In Zoom Out Reset image size

Radial velocity measurements with the VLT-UT1 CRyogenic high-resolution infraRed Echelle Spectrograph (CRIRES) presented in L. Malo et al. (in preparation) confirm the star's multiplicity, and yield radial velocities of 14.69 ± 0.53 km s⁻¹ and 14.43 ± 0.75 km s⁻¹, for A and B, respectively. The vsin i measurements (A: 15.52 ± 2.01 km s⁻¹; B: 20.40 ± 2.58 km s⁻¹) are also indicative of youth, where Reiners et al. (2009) considers vsin i > 20 km s⁻¹ the minimal condition for a "fast rotator."

While it is outside the normal spatial bounds (Torres et al. 2008) of the TW Hydra association (its R.A. is slightly lower than that of TWA 21, at 10^h13^m), its UVW kinematics are consistent with TW Hydra, as is its deblended isochrone position (Figure 4). Its gravity measurements (Figure 5) indicate it is extremely young.

TWA 8AB was listed by Jao et al. (2003) as RXJ1132-264AB and was not included in the TW Hydra analysis of Weinberger et al. (2013); our information is consistent with their conclusions based on other members of TW Hydra in that analysis.

G 165-8AB appears to be between the ages of β Pictoris and AB Doradus based on its gravity measurement and deblended H-R diagram positions. There are two highly discrepant radial velocities published for this system, +8 ± 0.1 km s⁻¹ (Montes et al. 2001) and −7.5 ± 6.5 km s⁻¹ (Gizis et al. 2002). We would normally choose the former value due to its higher precision, but the kinematics derived using that value agree with no known association. Ignoring radial velocities altogether, we find best-fit matches to the Tucana-Horologium (γ = 0.58, best-fit RV −6.2 km s⁻¹), Carina (γ = 0.62, best-fit RV −10.4 km s⁻¹), and Columba (γ = 2.40, best-fit RV −12.5 km s⁻¹) associations, all with estimated ages of 30 Myr. It seems likely that the Gizis et al. (2002) radial velocity is accurate, and the uncertainty is accounting for orbital motion of the binary. Using the latter radial velocity, the best match is to Carina, though BANYAN favors Columba.

Unusually, this system is a northern hemisphere target well outside the published spatial boundaries of all three associations. This suggests that either there is a fourth 30 Myr old association nearby, possibly containing two other northern hemisphere objects thought to be members of the 30 Myr old Columba association: HR 8799 (Marois et al. 2008; Baines et al. 2012) and κ And (Carson et al. 2013), or that the existing assumptions about the boundaries of the known associations are incorrect.

SCR 1425-4113AB is a curious system. High-resolution ESPaDOnS spectra of this target show that it is a spectroscopic binary, and an extremely rapid rotator (vsin i = 95.0 ± 8.1 km s⁻¹), with a radial velocity of (−1.2 ± 1.3 km s⁻¹). Its lithium EW was also measured at 595 ± 20 mÅ, confirming its youth. The deblended magnitudes of both of its components place them above (though consistent with) the TW Hydra isochrone on an H-R diagram, and the gravity measurements (which should not be affected by multiplicity) place it at lower surface gravity than TWA 8AB and SCR 1012-3124AB (though still within the errors).

Our kinematics find it to be a much better fit to β Pictoris (γ = 0.59) than TW Hydra (γ = 6.46), but BANYAN finds 75% probability of it being a TW Hydra member, compared with 24% probability that it is actually a β Pictoris member. It is additionally outside the normal spatial boundaries of TW Hydra (being a full hour of R.A. "higher" than the highest current member, TWA 18, at 13^h21^m R.A.). We nevertheless consider it as a potential TW Hydra member because with its extremely high luminosity, low surface gravity, and position relatively close to TW Hydra, it either must be an outlying member of TW Hydra or something equally young.

SCR 2010-2801AB was found to be a binary by Bergfors et al. (2010). The system was observed with HST-FGS on 2009 April 26, as seen in Figure 19. The resulting separation agrees with the separation published in Bergfors et al. (2010).

Zoom In Zoom Out Reset image size

LEHPM2-0783 has X-ray emission and the strongest Hα emission in the sample, despite its cool temperatures. It lies above the main sequence, though it is too red to be seen on Figure 4. It is, however, not a potential kinematic match to any known young association. It may be younger than a typical field star, but we have no reason to suspect it is anything more than an unresolved binary.

L 755-19 is kinematically consistent with both Castor and Argus. It lies within the upper main-sequence locus, above the AB Doradus isochrone, which would be more consistent with the Argus association (younger than AB Dor) or a main-sequence binary than the Castor moving group (older than AB Dor). L 755-19 is a hotter star than AP Col (Riedel et al. 2011), and its surface gravity is comparable with either field stars or Argus-age objects. A radial velocity would go a long way toward determining the status of this star, given that the best-fit Argus RV is −25.4 km s⁻¹, and the best-fit Castor RV is −9.3 km s⁻¹, and a Castor/main-sequence assignment would imply it is a binary.

SCR 2033-2556 is consistent with being a member of β Pictoris, but its high luminosity suggests it is an unresolved binary. Observations with ESPaDOnS (L. Malo et al., in preparation) suggest a lithium EW of 510 mÅ, confirming the star's youth.

GJ 803 (A)/GJ 799AB (BC) is one of the nearest young systems (see Table 10) to the Sun, and a prototypical member of the β Pic association (Barrado y Navascués et al. 1999). AU Mic (unobserved by CTIOPI) is known to have a dust disk, and several authors note that the AU-AT Mic separation is very large (at least 0.23 pc) and "must be very fragile and will soon be torn apart by third bodies" (Caballero 2009). This system is a well-established member of the 10 pc sample with parallaxes from Hipparcos and the Yale Parallax Catalog (van Altena et al. 1995); our results agree with the published values.

In our first epoch (2003 July 9) GJ 799 A and B were separated by 28 @ 171°; in our final epoch (2012 July 31, Figure 20) they were separated by 23 @ 153°.

Zoom In Zoom Out Reset image size

GJ 1284AB is identified as a potential member of Columba by kinematic analysis; however, Torres et al. (2006) identified it as a spectroscopic binary, which implies it is a system of main-sequence stars, and the kinematic match is spurious.