A SPECTROSCOPIC SURVEY OF MASSIVE STARS IN M31 AND M33^*

Most of the 1895 stars that we classify here were selected from the LGGS based upon their colors and magnitudes. The primary goal was to isolate the OB population, but because there is some spread in the average reddening (see Massey et al. 1986, 1995a, 2007b) we relied heavily on the Johnson reddening-free index $Q=(U-B)-0.72(B-V)$. Our photometric criteria were designed for obtaining spectra of OB stars with absolute visual magnitudes brighter than ${M}_{V}=-5.5$, roughly corresponding to 50 M_⊙ for the youngest O stars, and to 25 M_⊙ for older (5 Myr) O stars. Such stars will have V = 19.5 or brighter. The photometric criteria we adopted are given in Table 1. Generally for the fainter stars (19.5 ≥ V > 18.0) we insisted on a reddening-free index corresponding to an effective temperature of ∼23,000 K or hotter (see, e.g., Table 2 in Massey 1998c based upon Kurucz 1979, 1992), i.e., roughly a spectral type of B2.5 or earlier (see Table 3 in Massey 1998b). However, we imposed the additional restriction that both U− B and B − V colors must be reasonable as well, as shown by the other photometric criteria. For the brighter stars, we relaxed the requirements on Q as there were fewer stars, and identifying B- and A-type supergiants would be useful. In order to avoid the plethora of foreground dwarfs that would contaminate the latter sample by mid-A- or early-F-type, we imposed a U − B cut-off as a function of B − V.

Table 1. Photometric Criteria Used

Parameter	Criterion
	V ≤ 18.0	$19.5\geqslant V\gt 18.0$
B − V	≤0.4	≤0.4
U − B	≤0.05−1.125(B − V)	≤−0.4
Q	≤0.0	≤−0.7

Download table as:  ASCIITypeset image

The effect of these selection criteria are shown in Figures 3 and 4. The black points are the data from the LGGS, with the only restriction that V ≤ 19.5. Stars which meet the photometric criteria in Table 1 are shown by the magenta filled circles. The majority of such stars are found in the upper left region, bound by $B-V=0.4$ (vertical cyan line) and $Q\lt 0.7$ (diagonal cyan line). For the brighter stars (V ≤ 18.0) we also included potential B- and A-type supergiants; these are the stars found in the region bounded by the two blue lines, representing Q ≤ 0.0 and $U-B\leqslant 0.05-1.125(B-V)$. The latter criterion was imposed to avoid the large number of foreground stars, as demonstrated by the right panels in each figure. Note that there is very little, if any, expected contamination after applying our photometric criteria.

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Once the stars were selected photometrically, an additional "isolation" criterion was imposed. We considered the contaminating flux from all other neighboring stars (using the complete LGGS catalog) assuming 15 seeing. We required that the contamination be less than 20% at B, although stars with contamination between 10% and 20% were assigned a lower priority in the assignments. For M31 this took our original photometrically selected sample of 1163 stars and reduced it to 517, i.e., only 44% of the stars were sufficiently isolated. For M33, this took our photometrically selected sample of 1867 stars and reduced it to 619 (i.e., only 33% were sufficiently isolated).⁶

In addition to the primary program, we also included a number of other targets of potential interest, but at lower priority in assigning fibers. These include the M31 and M33 confirmed (and suspected) YSGs (ranks 1 and 2) from Drout et al. (2009) and Drout et al. (2012), along with other previously unobserved YSG candidates selected as described in those papers. The latter we expect to be heavily dominated by foreground dwarfs. This brought the total potential stars to observe to 3885 for M31, and 1785 for M33.

All of the observations on which our new spectroscopy is based come from the 6.5 m MMT telescope used with the 300 fiber positioner Hectospec (Fabricant et al. 2005).⁷ This instrument provides a 1° field-of-view, with a choice of two gratings: (1) a 270 line mm⁻¹ blazed at 5000 Å, which covers the entire visible region (3650–9200 Å) at once with a resolution of ∼5 Å, and (2) a 600 line mm⁻¹ grating blazed 6000 Å, which covers 2300 Å at one time at a resolution of ∼2 Å. The fibers have a diameter of 15. The majority of the new observations reported here were obtained with the higher dispersion grating nominally centered at 4800 Å but actually covering 3700–6000 Å. Some additional new spectral types are those of the non-WRs found by Neugent & Massey (2011) and Neugent et al. (2012a); these data were obtained with the lower-dispersion grating, with details reported in that paper.

Observations with Hectospec are made in a unique "self-staffed" queue mode. The observers are present for specific nights corresponding to the amount of time assigned by the telescope allocation committee, but they obtain observations guided by the schedule provided by the queue manager. The schedule is juggled so that at least some of the objects are the observer's own. The great advantage of this scheme is that weather losses are distributed among all of the participants. The first year of observations were obtained during the Fall of 2009, with nominally three nights assigned (2009B-0149), but poor weather during the semester meant we obtained only four fields in M31 and one field in M33. Each observation consisted of 3 × 45 minute exposures. The second year (Fall 2010), two nights were nominally assigned (NOAO 2010B-0260), allowing us to observe three additional fields in M33, two with 3 × 45 minute exposures and one with 3 × 40 minute exposures.

The spectra were all run through the standard SAO/CfA Hectospec IRAF⁸ pipeline (Mink et al. 2007) by Susan Tokarz of the OIR Telescope Data Center; an overview of the process was given by Drout et al. (2009). Note that SAO Telescope Data Center has just reprocessed all of their past Hectospec data, and performed a public release in late 2015; we have given permission for the data contained in the present paper to be included in this archive, and so as to avoid multiple versions floating around the astronomical community, we are not making our spectra separately available here.

The new spectral types were classified on the standard MK system with reference to Walborn & Fitzpatrick (1990), Jaschek & Jaschek (1990), Morgan et al. (1978), and Jacoby et al. (1984). The initial pass was done by B.M.S., with a careful reality check by K.F.N. Although the basic tenets of stellar spectral classification are well described elsewhere, we briefly summarize our procedure here in case it proves useful to others.

For the initial classification into spectral class, we followed this algorithm. If TiO molecular bands were present, the star was of M-type. (There were not many of these among our stars of course, given our photometric selection criteria; they were mainly introduced into the observing sample by the attempt to be inclusive for YSGs.) Otherwise, the Balmer lines Hδ, Hγ, and Hβ were used as preliminary indicators of spectral type. If they were dominant, the star is early-type (O, B, or A). If that was true, next we checked for the presence of He i and/or He ii. If neither were present, the star was of A-type. If He i was present, but not He ii, then the star was of B-type. If He ii was present, then the star was of O-type. If instead the Balmer lines were relatively weak, the G band at λ4308 was used as the primary discriminant. If the G band was completely absent (and the Balmer lines weak), the star was dismissed as an F-type foreground dwarf. If the G band was present, but weaker than Hγ, then the star was considered of F-type. If it was comparable to (or stronger) than Hγ, then the star was of G-K-type. If so, we next compared Hγ to Fe i λ4325. If Hγ was equal to or stronger than Fe i, then it was a G star; if not, then it was of K-type. We then classified the spectrum and assigned a luminosity type (and hence membership) as described below. To briefly summarize:

Of course, in addition to the stars with classical MK spectral types, our spectroscopy identified candidate luminous blue variable (cLBV) stars and even a new WR as described in more detail below.

In Table 2, we list the number of spectroscopically confirmed members of M31 and M33 as a function of type, denoting the number that are classified here for the first time. There is some overlap between the categories, i.e., a star might have an uncertain spectral type of A9-F0 I, in which case it will be counted in each category. Alternatively, a star might be counted twice because its spectral type is composite, i.e., WR+O. Nevertheless, these numbers are representative of the state of our knowledge of the resolved stellar populations of our closest spiral neighbors.

Many of the stars shown by Drout et al. (2009) and Drout et al. (2012) to be YSG members of M31 and M33 have subsequently been reclassified with MK types in this paper, and thus although no new "generic" YSGs are listed in Table 2, that is somewhat deceptive as these stars now are included under the more exact F I or G I headings. Similarly, many of the RSGs identified by Massey et al. (2009a) were also classified in that paper. We discuss each of these groups below.

cLBVs: Hubble & Sandage (1953) identified several blue, luminous stars in M31 and M33 that displayed spectacular photometric variability during a 40 year period; these stars, along with the Galactic and MC prototypes P Cygni, S Doradus, and η Car, are now called LBVs, following the suggestion of Conti (1984). Their spectra are variable, with a particular LBV sometimes showing an emission-line spectrum (either Ofpe/WN9 or dominated by permitted and forbidden Fe ii line emission along with He ii emission, or a P Cygni like spectra), or in a cooler state due either to a larger radius induced by the star's proximity to the Eddington limit (Groh et al. 2011), or the formation of a pseudo-photosphere in the stellar wind, leading to an F-supergiant-like spectrum (see, e.g., Massey 2000). Over the years, a number of other stars have been identified in M31 and M33 that have spectra indistinguishable from those of the classical LBVs. Some were observed spectroscopically because the stars were known to have emission as shown by narrow-band Hα imaging (e.g., King et al. 1998; Massey et al. 2007a), or due to UV-excess (Massey et al. 1996), or simply discovered "accidentally" as part of spectroscopic surveys, such as that of Massey (2006) and the present paper. Many of these "candidate" LBVs have now been shown to have photometric and/or spectral variability characteristic of true LBVs (Clark et al. 2012; Lee et al. 2014; Humphreys et al. 2015; Sholukhova et al. 2015). It is not clear what these cLBVs would have to do to be sanctified as true LBVs; the criteria needed seem to have been applied inconsistently and somewhat idiosyncratically within the literature. For instance, Humphreys et al. (2015) describes J004526.62+415006.3 as only the fifth "confirmed" LBV in M31, despite the fact that other candidate LBVs have shown similar variability, such as J004051.59+403303.0 (Sholukhova et al. 2015).

Our current spectroscopy has identified nine newly found cLBVs in these two galaxies. These new additions bring the number of known or suspected LBVs to 24 for M31 and 29 to M33.⁹ In Figure 5, we illustrate the spectra of the newly found cLBVs, and we list all of the known or suspected LBVs in M31 and M33 in Tables 3 and 4. With one exception, we exclude the so-called Ofpe/WN9 stars, now more commonly referred to as WN9-11 WRs (Crowther & Smith 1997). We realize that these late-type WNs stars have a strong linkage to LBVs. In fact, Neugent & Massey (2011) classified the star J013418.37+303837.0 as "Ofpe/WN9" without realizing that this star is actually M33 Var 2, one of the early prototypes of the LBV class (Humphreys 1978). The lack of good coordinates for the confirmed LBVs is part of the reason for this oversight; a problem which we believe we have remedied here. To the best of our knowledge, the Var 2 has never before shown a late-type WN spectrum; Humphreys (1978) describes the star has having an A-F-type spectrum along with Fe ii and [Fe ii]. Similarly, the star J013410.93+303437.6 was described as a newly found candidate LBV by Massey et al. (2007a), when in fact this is the Hubble & Sandage (1953) star M33 Var 83.

Zoom In Zoom Out Reset image size

The star J0133395.2+304540.5 was classified as "B0.5Ia+WNE" by Neugent & Massey (2011), but Clark et al. (2012) has demonstrated strong spectral variability, and we remove it here from the list of WRs, and include it now as an LBV candidate.

The star J013429.64+303732.1 was described by Massey et al. (2007a) as being a peculiar B8 I spectrum, showing narrow emission superimposed on a very broad component. In Figure 6 we show what appears to be changes that have taken place in the Hβ profile over the three years after the Massey et al. (2007a) spectrum was taken: the broad emission component is unchanged, but there is now an absorption component where there had been emission. We cannot rule out an instrumental effect: the earlier spectrum was taken with a 3'' diameter fiber with Hydra, while the later spectrum was taken with a 15 fiber with Hectospec, each with different sky placements, so it is possible that nebular contamination is responsible for this change. In addition, it is a minor change compared to the spectral variability found for true LBVs, and we leave its description as a "candidate."

Zoom In Zoom Out Reset image size

WR stars: Neugent et al. (2012a) and Neugent & Massey (2011) published the results of comprehensive surveys for WRs in M31 and M33, respectively. Since that time Shara et al. (2016) found one additional WR star in M31.¹⁰ In addition, Neugent & Massey (2014) spectroscopically confirmed five of the previously unobserved M33 WR candidates listed by Neugent & Massey (2011). (The sixth proved to be an Of-type star.) In the current paper we have identified one additional WR star in M33, J013344.05+304127.5, a B0 I+WN pair whose spectrum is dominated by the B0 I component. We show the spectrum of this star in Figure 7. This brings the total number of WRs identified in M33 to 211, the 206 cataloged by Neugent & Massey (2011), minus J013395.2+304540.5 (which we now call a cLBV as discussed above), plus the addition of the five WRs confirmed by Neugent & Massey (2014), plus the one new one here. All of the new additions are of WN-type.

Zoom In Zoom Out Reset image size

O-type stars: all O-type stars in our selected magnitude range will be M31/M33 members, even the dwarfs. The classification of O-type spectra is quite straightforward, relying upon the relative strengths of He i and He ii. The presence of He ii λ4686 emission and/or N iiiλλ4634,42 is then the primary luminosity indicator; more subtle aspects of this are described in both Walborn & Fitzpatrick (1990) and Sota et al. (2011).

As Table 2 shows, most of the O-type stars known in M31 (83%) and M33 (60%) were identified as a result of our new spectroscopy here. One of the stated goals of the current project was to improve our knowledge of the unevolved massive star populations of these galaxies, and while we have made an improvement, the number of O-type stars now classified in M31 (64) and in M33 (130) are only a few percent of the total number of expected O-type stars in these galaxies. Current evolutionary models for solar metallicity predict there should be roughly 15× more O-type stars than WRs (see Table 3 in Georgy et al. 2012). Thus, we expect on the order 2000–3000 O-type stars in each of these galaxies. So, while we have made a substantial improvement, there is still a way to go in terms of characterizing the unevolved massive star populations of these two galaxies.¹¹

In Figure 8, we show examples of several of the newly found O-type stars in this study; the stars were picked as nice examples of early and late O-type dwarfs and supergiants. We note that the O4 V star J004258.29+414645.9 is the earliest type star known in M31, while the O5 If star J013413.06+305230.0 is the earliest known supergiant in M33. These spectra are remarkable mostly for their being so ordinary: they are typical of any similar stars in the Milky Way or MCs. The only extraordinary fact is that these stars are ∼10× more distant than similar objects in the Large MC or SMC.

Zoom In Zoom Out Reset image size

B-type supergiants: we expect only B-type supergiants and early B-type giants to be M31 members; at the same time, we expect no B-type foreground dwarfs to contaminate the sample. The classification of B stars relies upon the relative strengths of Si ivλ4089, Si iiiλ4553, and Si ii λ4128, with the relative strengths of Mg iiλ4481 and He iλ4471 being a secondary indicator for B stars later than B2. The luminosity criteria for the earliest B stars relies upon a comparison of Si iv λ4116 to He i λ4121, or Si iii λ4553 to He i λ4387. At later B-types the strengths and shape of the Balmer lines were used to confirm a supergiant classification.

Our study here is responsible for the vast majority (86% for M31, and 79% for M33) of the B-type supergiants spectroscopically identified in M31 and M33, respectively (Table 2). We illustrate a few examples in Figure 9.¹²

Zoom In Zoom Out Reset image size

A-type stars: we expect some later-type A stars to be foreground dwarfs; fortunately the strengths of the Fe ii and other metal lines, and the shape and strength of the Balmer lines, readily distinguish between the A-type supergiant members of M31/M33 and foreground dwarfs. The spectral subtypes are straightforwardly determined by comparing the strengths of the Ca ii H (λ3968) and K (λ3934) lines, as the former is blended with H and the strengths of the Balmer lines decrease as one goes from A0 to A9; an increase in the strengths of the metal lines also was taken into account, as these reach a maximum at A5.

We see from Table 2 that about 90% of the A-type supergiants known in M31 and M33 have been identified as a result of our new study. We show some examples of A-type spectra in Figure 10.

Zoom In Zoom Out Reset image size

F-type supergiants: here we expect foreground dwarfs to dominate the sample (see, e.g., Figures 3 and 4). However, the F-type supergiants are very recognizable, as they have very strong metal lines. In addition, for F0-4, supergiants will have strong Ti iiλ4444 relative to Mg iiλ4481. From F5-9, supergiants will not show the G band, and Sr iiλ4077 will have about half the strength of Hδ. The spectral subtypes themselves were determined primarily by comparing the Ca i λ4226 line with Hγ.

Most of the F-type supergiants known in these galaxies were discovered by the radial velocity surveys by Drout et al. (2009) and Drout et al. (2012), and are classified here for the first time; most of the others were recently classified by Humphreys et al. (2013). We show examples of our spectra in Figure 11.

Zoom In Zoom Out Reset image size

G-type supergiants: again, foreground dwarfs will completely dominate the sample. The occasional supergiant is easily recognized as having very sharp lines with the components of the G band clearly separated. In addition, supergiants have Sr iiλ4216 comparable to that of Ca iλ4226. The spectral subtypes were determined by comparing Hγ and Fe iλ4325, and Hδ with Fe iλ4045, with stronger Balmer lines indicating earlier types.

All of the known G-type supergiants in M31 and M33 have been classified as such as part of the current paper (Table 2). All but two of the M31 stars had been recognized as YSGs by the radial velocity surveys of Drout et al. (2009) and Drout et al. (2012). We show examples in Figure 12.

Zoom In Zoom Out Reset image size

In Figures 13 and 14, we show the extent that our spectroscopy has affected the interpretation of the CMDs. On the left side we show the upper portion of the CMDs from Figures 1 and 2. On the right, we have used the spectroscopic information now available to eliminate stars known to be foreground stars. We see that the thick band of yellow stars has become thinner. In addition, we indicate in red the stars for which our spectroscopy has established membership. Primarily these are stars in the "blue plume," but with a smattering of YSGs and RSGs. The contrast in how easily seen the red points are (denoting the stars with spectroscopy) demonstrates the fact that this work is considerably more advanced in M33 than in M31 right now (see also Table 2), a situation we hope to rectify in the coming years.

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Of course, we were most intrigued to see where these stars fell in a physical HRD, with comparison to evolutionary tracks. Note, however, that there are strong selection effects: the numbers of WR stars are nearly complete (Neugent et al. 2012a; Neugent & Massey 2011), the numbers of YSGs are probably complete (Drout et al. 2009, 2012), and the number of RSGs is complete for M33 (Drout et al. 2012) but not for M31, although work is in progress (Evans & Massey 2015). At the same time, we have argued that the number of O-type stars known in these systems is woefully incomplete, and that our current knowledge represents only a few percent of the total. So, care must be exercised in not over interpreting these data. Still, it is of obvious interest to see where these fall in the HRD.

We have chosen to then plot only stars of spectral types O–F (and/or classified as YSGs), ignoring then the cLBVs, LBVs, the WRs, and the RSGs. The argument is that detailed modeling work is needed for even approximate temperatures for the (c)LBVs and WRs, while the locations of RSGs require either model fitting or, at the least, K-band photometry, and all of this is best handled elsewhere. We have also ignored stars that were obviously composite (double-lined or otherwise composite).

For the O-G stars with MK types, we assigned effective temperatures based on their spectral types, using the effective temperature scale of Massey et al. (2005) for the O stars and Cox (2000) for B2 I through G8 I. (The B0-B2 I values were adjusted for a smooth transition.) Photometry was corrected for reddening assuming $E(B-V)=0.13$ for M31 and $E(B-V)=0.12$ for M33 (Massey et al. 2007b); these values are averages, but the scatter for early-type stars is fairly small. The extinction at V was assumed then to be ${A}_{V}=3.1E(B-V)$. For the stars without MK types (such as those classified as YSGs) we used the de-reddened photometry and the relation

$$\begin{eqnarray*}\mathrm{log}{T}_{{\rm{eff}}} & = & 4.008-0.915{(B-V)}_{0}\\ & & +0.9815(B-V{)}_{0}^{2}-0.36212(B-V{)}_{0}^{3},\end{eqnarray*}$$

which is derived from a combination of model atmospheres (see Drout et al. 2012) and empirical data (Cox 2000 and references therein), and is useful for supergiants over the range $3.6\lt \mathrm{log}{T}_{{\rm{eff}}}\lt 4.2$ and $-0.1\lt {(B-V)}_{0}\lt 1.6$.¹³ In M33 there were a few stars with "generic" O-type designations that we also ignored. The bolometric corrections were calculated using

$$\begin{eqnarray*}{BC} & = & -1131.425+827.397\times \mathrm{log}{T}_{{\rm{eff}}}-200.6348\\ & & \times {(\mathrm{log}{T}_{{\rm{eff}}})}^{2}+16.11878\times {(\mathrm{log}{T}_{{\rm{eff}}})}^{3}\end{eqnarray*}$$

based upon the values in Massey et al. (2005) and Cox (2000). Distances of 760 and 830 kpc (van den Bergh 2000) were assumed for M31 and M33, respectively.

We show the resulting HRDs in Figure 15. We have included the solar metallicity evolutionary tracks from Ekström et al. (2012); these are computed for solar metallicity (Z = 0.014) and thus are too low for M31 and too high for M33, but they give a reasonable idea of the mass ranges of the stars plotted.

Zoom In Zoom Out Reset image size

The two highest mass stars in our sample for M31 are both early-type B supergiants, J004422.84+420433.1 (B0.5 I) and J004057.86+410312.4 (B1-1.5 I). These stars have masses of ∼60 M_⊙ according to the evolutionary tracks. While there are doubtless many O stars this massive or more so, they will be considerably fainter than these supergiants, owing to the larger (more negative) bolometric correction; it is for this reason that the younger massive stars are so disproportionately missing in our spectroscopic sample, a point that has been emphasized many previous times (see, e.g., Massey et al. 1995b; Massey 1998c): an "old" (4 Myr) 40 M_⊙ star will be a B0 I with ${M}_{V}=-6.6$, while the same star when it was younger (0–1 Myr) will be an O6 V that is two magnitudes fainter visually! (See Table 1 in Massey 1998c.) The majority of the unevolved population we have identified in M31 is 20–40 M_⊙.

The HRD of M33 is very similar. Again, the highest luminosity (mass) star appears to are two early-type B supergiants, J013440.41+304601.4 (B1.5 I) and J013359.98+303354.7 (B1 Ie), and again the vast majority of the relatively unevolved massive stars in our sample are in the 20–40 M_⊙ range.

In Tables 5 and 6, we provide now the revised versions of the LGGS UBVRI catalogs of stars in M31 and M33. The photometry is identical to that of Massey et al. (2006) with the U – B corrections as noted in the subsequent erratum (Massey et al. 2011). We have included not only the new spectral classifications derived here, but also updated the tables with all of the spectral types and information that have been published in the intervening years. We have also included some other new information to make these tables more useful, including membership criterion based upon the spectroscopy, an index indicating the degree of crowding, and a value for the galactocentric distance of each star within the plane of its galaxy.

Table 5. Updated M31 LGGS Catalog

Star	ρa	α2000	δ2000	V	σV	B − V	${\sigma }_{B-V}$	U − Bb	${\sigma }_{U-B}$ b	V − R	${\sigma }_{V-R}$	R − Ib	${\sigma }_{R-I}$ b	NV	NB	NU	NR	NI	Cwdc	Memd	Referencese	Type
J003701.92+401233.2	1.24	00 37 01.91	+40 12 33.1	19.862	0.017	−0.021	0.021	−0.928	0.015	0.204	0.023	99.999	99.999	1	2	1	1	0	I	⋯	⋯	⋯
J003701.93+401218.4	1.24	00 37 01.92	+40 12 18.3	18.739	0.008	1.494	0.015	0.945	0.036	0.946	0.014	99.999	99.999	1	2	1	1	0	I	⋯	⋯	⋯
J003702.03+401141.4	1.23	00 37 02.02	+40 11 41.3	21.225	0.043	1.362	0.085	99.999	99.999	0.748	0.049	0.694	0.024	1	1	0	1	1	I	⋯	⋯	⋯
J003702.05+400633.5	1.18	00 37 02.04	+40 06 33.4	21.091	0.044	0.050	0.061	−1.110	0.052	0.042	0.074	99.999	99.999	1	2	1	1	0	I	⋯	⋯	⋯
J003702.13+400945.6	1.21	00 37 02.12	+40 09 45.5	16.091	0.006	1.287	0.007	0.983	0.007	0.792	0.010	99.999	99.999	1	2	1	1	0	I	0	1	fgd
J003702.24+401225.7	1.24	00 37 02.23	+40 12 25.6	20.765	0.029	1.584	0.072	99.999	99.999	1.030	0.036	1.262	0.021	1	2	0	1	1	I	⋯	⋯	⋯
J003702.38+400529.5	1.18	00 37 02.37	+40 05 29.4	22.427	0.165	1.359	0.498	99.999	99.999	0.574	0.180	0.892	0.072	1	1	0	1	1	I	⋯	⋯	⋯
J003702.44+400723.2	1.19	00 37 02.43	+40 07 23.1	22.461	0.173	−0.153	0.202	99.999	99.999	−0.018	0.252	99.999	99.999	1	2	0	1	0	I	⋯	⋯	⋯
J003702.47+401742.5	1.30	00 37 02.46	+40 17 42.4	18.026	0.007	1.324	0.011	1.150	0.023	0.732	0.011	0.688	0.008	1	2	1	1	1	I	0	1	fgd
J003702.51+401654.5	1.29	00 37 02.50	+40 16 54.4	18.768	0.008	0.687	0.012	−0.040	0.014	0.393	0.011	0.390	0.008	1	2	1	1	1	I	⋯	⋯	⋯

Notes. Table 5 is published in its entirety in the electronic edition. A portion is shown here for guidance regarding its form and content.

^aDeprojected distance within the plane of M31 in units of the Holmberg radius. Assumes a Holmberg radius of 953, inclination 770, and position angle of major axis of 350. At a distance of 760 kpc, a value of 1.00 for ρ corresponds to 21.07 kpc. ^b99.999 = no value. ^cCrowding: I—isolated (contamination <5%), S—somewhat isolated (5%–20%), C—crowded (20%–80%), X—extremely crowded (>80%). ^dMembership: 0—non-member, 1—member, 2—possible member, 3—unknown. ^eReferences for spectral type: (1) Drout et al. 2009; (2) current paper; (3) Massey et al. 2006; (4) Trundle et al. 2002; (5) Humphreys et al. 1988; (6) Humphreys 1979; (7) Neugent et al. 2012a; (8) Massey et al. 2009a; (9) Cordiner et al. 2011; (10) Massey 1998a; (11) Humphreys et al. 1990; (12) Cordiner et al. 2008a; (13) Massey et al. 1995a; (14) Massey et al. 2007a; (15) Humphreys et al. 2013; (16) Massey & Johnson 1998); (17) Hubble & Sandage 1953; (18) King et al. 1998; (19) Massey 2006; (20) Massey & Evans 2016.

Only a portion of this table is shown here to demonstrate its form and content. A machine-readable version of the full table is available.

Download table as:  DataTypeset image

Table 6. Updated M33 LGGS Catalog

Star	ρa	α2000	δ2000	V	σV	B − V	${\sigma }_{B-V}$	U − Bb	${\sigma }_{U-B}$ b	V − R	${\sigma }_{V-R}$	R − Ib	${\sigma }_{R-I}$ b	NV	NB	NU	NR	NI	Cwdc	Memd	Referencese	Type
J013146.18+302931.4	1.37	01 31 46.15	+30 29 31.3	21.027	0.061	0.090	0.113	0.266	0.118	1.012	0.111	99.999	99.999	1	1	1	1	0	X	⋯	⋯	⋯
J013146.18+302932.4	1.37	01 31 46.15	+30 29 32.3	20.560	0.068	0.645	0.117	99.999	99.999	0.564	0.115	99.999	99.999	1	1	0	1	0	C	⋯	⋯	⋯
J013146.20+302706.2	1.36	01 31 46.17	+30 27 06.1	21.057	0.032	1.857	0.084	99.999	99.999	0.924	0.036	99.999	99.999	1	1	0	1	0	I	⋯	⋯	⋯
J013146.21+302026.9	1.36	01 31 46.18	+30 20 26.8	21.179	0.038	0.962	0.066	0.749	0.096	0.588	0.047	99.999	99.999	1	1	1	1	0	I	⋯	⋯	⋯
J013146.25+301849.7	1.36	01 31 46.22	+30 18 49.6	16.359	0.005	0.811	0.007	0.470	0.007	0.460	0.007	99.999	99.999	1	1	1	1	0	I	0	1	fgd
J013146.26+302931.5	1.37	01 31 46.23	+30 29 31.4	20.247	0.063	0.912	0.128	0.223	0.131	−0.209	0.112	99.999	99.999	1	1	1	1	0	C	⋯	⋯	⋯
J013146.43+302048.4	1.35	01 31 46.40	+30 20 48.3	21.990	0.074	0.460	0.102	−0.155	0.089	0.308	0.097	99.999	99.999	1	1	1	1	0	I	⋯	⋯	⋯
J013146.64+301756.1	1.36	01 31 46.61	+30 17 56.0	22.075	0.080	1.498	0.156	99.999	99.999	1.121	0.086	99.999	99.999	1	1	0	1	0	I	⋯	⋯	⋯
J013146.73+303118.0	1.37	01 31 46.70	+30 31 17.9	22.110	0.080	1.256	0.153	99.999	99.999	0.710	0.088	0.805	0.037	1	1	0	1	1	I	⋯	⋯	⋯
J013146.81+302614.2	1.35	01 31 46.78	+30 26 14.1	22.258	0.094	1.856	0.274	99.999	99.999	0.952	0.104	1.115	0.044	1	1	0	1	1	I	⋯	⋯	⋯

Notes. Table 5 is published in its entirety in the electronic edition. A portion is shown here for guidance regarding its form and content.

^aDeprojected distance within the plane of M33 in units of the Holmberg radius. Assumes a Holmberg radius of 308, inclination 560, and position angle of major axis of 230. At a distance of 830 kpc, a value of 1.00 for ρ corresponds to 7.44 kpc. ^b99.999 = no value. ^cCrowding: I—isolated (contamination <5%), S—somewhat isolated (5%–20%), C—crowded (20%–80%), X—extremely crowded (>80%). ^dMembership: 0—non-member, 1—member, 2—possible member, 3—unknown. ^eReferences for spectral type: (1) Drout et al. 2012; (2) current paper; (3) U et al. 2009; (4) Neugent & Massey 2011; (5) Humphreys et al. 2013; (6) Humphreys 1980; (7) Massey & Johnson 1998; (8) Massey et al. 2007a; (9) Massey et al. 1996; (10) Massey 1998a; (11) Massey et al. 1995a; (12) Massey et al. 2006; (13) Neugent & Massey 2014; (14) Drissen et al. 2008; (15) van den Bergh et al. 1975; (16) Abbott et al. 2004; (17) Monteverde et al. 1996; (18) Kehrig et al. 2011; (19) Hubble & Sandage 1953; (20) Cordiner et al. 2008b; (21) Oey & Massey 1994; (22) Valeev et al. 2009; (23) Clark et al. 2012; (24) Humphreys 1978; (25) Urbaneja et al. 2005.

Only a portion of this table is shown here to demonstrate its form and content. A machine-readable version of the full table is available.

Download table as:  DataTypeset image

We have slightly modified the coordinates to place them essentially on the UCAC3/4 system (Zacharias et al. 2013); the original LGGS coordinates were tied to the USNO-B1.0 catalog (Monet et al. 2003), requiring small (∼03) corrections to place the stars on the same system as guide stars suitable for observing with Hectospec. For M31:

$$\begin{eqnarray*}&&{\alpha }_{\mathrm{new}}={\alpha }_{\mathrm{old}}-{0.0065}^{{\rm{s}}}\end{eqnarray*}$$

$$\begin{eqnarray*}&&{\delta }_{\mathrm{new}}={\delta }_{\mathrm{old}}+2\buildrel{\prime\prime}\over{.} 283-3\buildrel{\prime\prime}\over{.} 346{\delta }_{\mathrm{old}},\end{eqnarray*}$$

where the δ_old in the last equation is expressed in radians. For M33, Neugent & Massey (2011) find simpler corrections:

$$\begin{eqnarray*}&&{\alpha }_{\mathrm{new}}={\alpha }_{\mathrm{old}}-{0.030}^{{\rm{s}}}\end{eqnarray*}$$

$$\begin{eqnarray*}&&{\delta }_{\mathrm{new}}={\delta }_{\mathrm{old}}-0\buildrel{\prime\prime}\over{.} 14.\end{eqnarray*}$$

The result of these corrections is that the LGGS designations no longer exactly correspond to the listed coordinates.

In addition, we have included an indication of how crowded based on the degree of contamination we expect in a 15 fiber or slit under 10 seeing conditions in either B, V, or R, whichever is worse.