Lkh 101 and the YOUNG CLUSTER in NGC 1579

LkH 101 AND THE YOUNG CLUSTER IN NGC 1579 G. H. Herbig, Sean M. Andrews, and S. E. Dahm Institute for Astronomy, University of Hawaii, 2680 Woodlawn Drive, Honolulu, HI 96822 Receivved 2004 May 26; accepted 2004 June 3

ABSTRACT

The central region of the dark cloud L1482 is illuminated by LkH 101, a heavily reddened (AV 10 mag) ; 3 high-luminosity (8 10 L) star having an unusual emission-line spectrum plus a featureless continuum. About 35 much fainter (mostly between R ¼ 16 and >21) H emitters have been found in the cloud. Their color- magnitude distribution suggests a median age of about 0.5 Myr, with considerable dispersion. There are also at least ﬁve bright B-type stars in the cloud, presumably of about the same age; none show the peculiarities expected of HAeBe stars. Dereddened, their apparent V magnitudes lead to a distance of about 700 pc. Radio observations suggest that the optical object LkH 101isinfactahotstarsurroundedbyasmallHii region, both inside an optically thick dust shell. The level of ionization inferred from the shape of the radio continuum corresponds to a Lyman continuum luminosity appropriate for an early B-type zero-age main-sequence star. The VI color is consistent with a heavily reddened star of that type. However, the optical spectrum does not conform to this expectation: the absorption lines of an OB star are not detected. Also, the [O iii] lines of an H ii region are absent, possibly because those upper levels are collisionally deexcited at high densities. There are several distinct contributors to the optical spectrum of LkH 101. The H emission line is very strong, with wings extending to about 1700 km s1, which could be produced by a thin overlying layer of hot electron scatterers. There is no sign of P Cygni type mass ejection. Lines of Si ii are narrower, while the many Fe ii lines are still narrower and are double with a splitting of about 20 km s1.Linesof[Feii], [O i], and [S ii] are similarly sharp but are single, at the same velocity as the Fe ii average. Work by Tuthill et al. allowed the inference, from K-band interferometry, that the central source is actually a small horseshoe-shaped arc about 0B05 (35 AU) across. A tipped annulus of that size in rotation about a 15 M star would produce double spectrum lines having about the splitting observed for Fe ii. The totality of observational evidence encourages the belief that LkH 101 is a massive star caught in an early evolutionary state. Key words: open clusters and associations: individual (NGC 1579) — stars: emission-line, Be — stars: formation — stars: individual (LkH 101) — stars: pre–main-sequence Online material: machine-readable table

1. INTRODUCTION since then been investigated in increasingly greater detail (Herbig 1971; Allen 1973; Thompson et al. 1976; Simon & Current belief is that the formation of a massive star takes Cassar 1984; Hamann & Persson 1989; and others). place deep in the parent cloud, behind very heavy extinction, Early radio observations (Brown et al. 1976; Cohen 1980; and hence will not be optically accessible (Palla & Stahler Purton et al. 1982) indicated that there were two principal 1990; Bernasconi & Maeder 1996). But when such a star contributors to the radio continuum at NGC 1579. The first is evolves further to the zero-age main sequence (ZAMS), its a point source at the position of LkH 101, believed to be a radiation and wind will clear out the neighborhood and the star hot star plus a very small H ii region, both behind an optically could at some time become optically detectable. We examine thick dust shell. Later, Hoare et al. (1994) and Hoare & here the possibility that LkH 101, in the nebula NGC 1579, Garrington (1995), from MERLIN interferometry, determined may be such a transition object. mean sizes for this source of 0B55 at 18.7 cm and 0B31 at 6 cm. NGC 1579 is a clump of bright nebulosity about 20 across, At 10 m, Danen et al. (1995) found it unresolved and set an lying in a dark cloud (L1482) north of the main Taurus-Auriga upper limit of 0B34 on its FWHM. This is the object to which cloud complex, a line of sight that also passes through the more the optical spectroscopy of LkH 101 refers. The second radio distant Per OB2 association. Although there are several stars continuum component is an extended H ii envelope, of di- illuminating their own small reflection nebulae nearby, in the ameter of the order of 10, the whole being embedded in a era of photography in blue-violet light there was no obvious clumpy H i cloud about 50 in diameter (Dewdney & Roger source of illumination of NGC 1579 itself, nor was any cluster 1986). of embedded stars apparent. In 1956 it was found (Herbig Sharpless (1959) probably included NGC 1579 as S222 in 1956) that at the edge of the bright nebulosity there is a very his catalog of H ii regions because it is so much brighter on the faint, very red star (V 16) having a powerful H emission red Palomar plates than on the blue. However, contrary to that line and an unusual emission-line spectrum. Subsequent work and to expectation from the radio results, the [O iii] kk4957, has shown that the spectra of star and nebula are identical (x 4) 5007 lines, characteristic of H ii regions, are not present in the and that the polarization of the nebula points to a source at that spectrum of the nebula or of LkH 101. Clearly, NGC 1579 is location (Redman et al. 1986), so this is the illuminating source a reflection nebula and owes its redness to interstellar extinc- of the nebula. The spectrum of this star, named LkH 101, has tion and to the nature of its illuminating source. But why that 1233 1234 HERBIG, ANDREWS, & DAHM Vol. 128

Fig. 1.—False-color image of NGC 1579 constructed from 300 s B-, V-, and R-band exposures. The frame is roughly 7A5 (1.5 pc for a distance of 700 pc) on a side, with north up and east to the left. source does not show the spectrum of a H ii region is another continuum is indeed that of a B0.5 ZAMS star, Cohen (1980) matter. found from the narrowband colors of LkH 101 that If the radio continuum spectrum of the core source in LkH AV ¼ 9:1 0:5 mag. He found also from recombination the- 101 is the free-free emission of a spherically symmetric H ii ory that the Balmer decrement, those lines assumed optically region, then the Lyman continuum (Lyc) flux required to thin, gave AV ¼ 12:5 1 mag. Subsequent investigators maintain that ionization can be obtained (Harris 1976; Brown (Thompson et al. 1976; McGregor et al. 1984; Rudy et al. et al. 1976; Knapp et al. 1976). All found, on the basis of the 1991; Kelly et al. 1994) using various procedures have found properties of OB stars tabulated by Panagia (1973), that the values of AV ranging from 9.7 to 15.8 mag. required Lyc flux could be supplied by a single star, if near If one simply assumes that AV ¼ 10 mag and a distance the ZAMS, of type B0 or B1.1 Assuming that the optical of 700 pc (obtained later in this paper), then the observed value of V ¼ 15:7leadstoMV ¼3:5. This is not incompatible with a normal B0.5 V, given the crudity of this cal- 1 More recent calculations (Vacca et al. 1996) based on later values of T e culation plus the scatter in the values of MV found in the and L and improved atmospheric models (Sternberg et al. 2003) predict sub- literature for that spectral type: Panagia (1973) gave 3.5 and stantially higher Lyc fluxes for OB stars, so a somewhat later B type for LkH 101 would follow. It is not possible to be more specific until such calculations Vacca et al. (1996) gave 4.1, while Andersen (1991) found are extended to types later than B0.5. 3.2 and 2.9 from two eclipsing binaries. The mass of a No. 3, 2004 NGC 1579 1235

TABLE 1 Optical and Near-Infrared Photometry of Stars in NGC 1579a

Star (4+) (35+) VB VV RV IRR IJ J H b H K SpT

1...... 29 54.12 15 19.8 ...... 17.61 1.74 ...... 2...... 29 55.03 15 11.8 ...... 17.57 1.04 0.59 ... 3...... 29 55.10 17 47.1 ...... 17.66 1.63 0.66 ... 4c ...... 29 55.56 15 02.4 ...... 17.49 1.34 0.80 ... 5...... 29 55.57 15 04.4 ...... (18.17) 1.25 ...

Notes.—Table 1 is presented in its entirety in the electronic edition of the Astronomical Journal. A portion is shown here for guidance regarding its form and content. Units of right ascension () are hours, minutes, and seconds, and units of declination ( ) are degrees, arcminutes, and arcseconds. a The BVRI colors have been converted to the Johnson-Kron-Cousins photometric system, whereas the JHK colors are on the new Mauna Kea ﬁlter system. b When only H and K measurements are available, this column contains the H magnitude in parentheses. c Integrated magnitudes of an unresolved binary.

B0.5 V according to the latter two sources is 19 and 13 M, (1979) found that between 1 and 160 m this ‘‘extended re- ; 4 respectively. gion’’ has a total L ¼ 1:2 10 L, still somewhat low. The With respect to total luminosity, the match with expectation explanation may simply be that most of the radiation of an early does not seem so satisfactory. The same authorities (above) B star is in the ultraviolet, and such integrations must miss that 4 give for L of a B0.5 V star (in units of 10 L): Panagia, 2.0; fraction that is not rethermalized by circumstellar dust or does Vacca et al., 6.2; and Andersen, 1.9 and 1.4. Actual integration not appear in the near-infrared free-free continuum. of the narrowband photometry of LkH 101 between 1.3 and Therefore, there is reason to suspect that LkH 101 may 25 m by Strecker & Ney (1974) and by Simon & Cassar indeed be a star of mass about 15 M in an interesting phase of (1984), corrected as above for extinction and distance, gives its early evolution. We proceed on that assumption. 0.8 in the same units. However, there is warm dust, presumably In what follows, we discuss (x 2) our optical and near- illuminated by the star, well away from the core. Harvey et al. infrared photometry of the heavily obscured cluster of stars

Fig. 2.—Internal photometric errors of the optical and infrared photometry. 1236 HERBIG, ANDREWS, & DAHM Vol. 128 surrounding LkH 101, (x 3) the optical spectrum of LkH TABLE 2 101 at high-resolution, and (xx 5, 6) the surrounding molecular H Emission Stars in NGC 1579 cloud and its contents. IH Table 1 W(H)a

2. THE STAR CLUSTER 187...... 10 8 2.1. Optical Photometry 188...... 27 22 189...... 30 79 Figure 1 is a false-color composite of NGC 1579 created 190...... 32 3 from 300 s B, V,andR images of this data set. Optical images 191...... 44 4 192...... 63 4 behind BVRC IC filters (hereafter we omit the subscripts on RC 193...... 70 9 and IC), covering an area of about 7A5 ; 7A5centeredonLkH 101, were obtained at the f/10 focus of the University of 194...... 72 8 Hawaii (UH) 2.2 m telescope in 1999 October. Conditions 195...... 78 54, 16 B 196...... 83 26 were photometric, the seeing FWHM about 0 7(inV ). Ex- 197...... 95 94, 32 posure times were 10, 60, and 300 s in each filter so both 198...... 100 61 the bright stars and the faint cluster population were mea- 199...... 105 20, 11 surable. The detector was a Tektronix 20482 CCD, the scale 200...... 107 6 0B22 pixel1. 201...... 111 69 All the data frames were corrected for bias by subtraction of 202...... 112 75, 35 a median-combined composite dark exposure. Any residual 203...... 118 4 bias was removed from individual frames by fitting a poly- 204...... 122 10 nomial to the overscan region and subtracting the fit across the 205...... 126 34 frame. Flat-field images were generated from observations of 206...... 132 21 207...... 139 224 the twilight sky. Quantum-efficiency variations (typically <1% 208...... 140 33 for this CCD) were removed by dividing the data frames by 209...... 151 51 these median-combined flat-field images. 210...... 157 372 Aperture photometry was performed using the DAOPHOT 211...... 180 6 package in IRAF. In the rare cases where a star fell on the 212...... 187 6, 7 boundary between bright nebulosity and a darker background, 213...... 192 8 the sky contribution was removed on an individual basis, not 214...... 194 1075 the usual annular procedure. Closely grouped stars were ana- 215...... 205 45 lyzed separately using point-spread function (PSF) fitting 216...... 215 24 photometry with the PSF+ALLSTAR tasks. The measured 217...... 225 3, 3 218...... 233 6 instrumental magnitudes were converted to the BVRI system 219...... 243 182, 127 using observations of Landolt (1992) standard stars over a 220...... 253 176, 102 range of air masses. The limiting magnitude is about V ¼ 22, 221...... 303 164 with completeness to V 20:5. Table 1 is a list of optical 222...... 304 20 magnitudes and colors for all stars detected above the 3 level a in at least two bandpasses. The coordinates were derived by All equivalent widths are in angstroms. Those with two measurements have the grism value first and reference to stars in the HST Guide Star and USNO-A Cata- the GMOS value second. logs. They are reliable at about the 100 level. The internal errors of the photometry are shown in Figure 2. The field was reimaged again with the same instrumentation to LkH 101. No H emitters were found in either of these on 2003 November 10, under photometric conditions but with fields. Equivalent widths of the H line in all the H emitters only average seeing and brightly moonlit sky. Good photom- (grism+GMOS detections) are given in Table 2. They are as- etry was possible for only the brighter stars in the field but was signed IH numbers in continuation of the numbering sys- sufficient to confirm the photometric zero points of the earlier tem of Herbig & Dahm. In Figure 3 all these stars are identi- observations. fied with their numbers in Table 1. The heavy and irregular extinction in NGC 1579 causes the There are very few H emitters, but many heavily reddened optical sample of Table 1 to be an incomplete representation of stars, east of LkH 101, probably because of greater fore- the faint cluster population. Nearly 100 stars that were detected ground extinction in that area, as suggested by the CO contours only in I are not listed there. in Barsony et al. (1990; their Figs. 11 and 12). Slitless grism spectra (6300–6700 8) covering the same area 2.2. Near-Infrared Photometry were obtained at the UH 2.2 m telescope in 1990, 1998, and 2003. The detectors, scale, etc., were as described by Herbig & Near-infrared JHK images of NGC 1579 were obtained in Dahm (2002) in an investigation of IC 5146 with the same 2002 August with the QUIRC infrared array (Hodapp et al. equipment. On these exposures about 19 faint stars having H 1995) at the f/10 focus of the UH 2.2 m telescope. The typical in emission were found (16 others were later found with GMOS seeing FWHM was 0B5(inK ). In this configuration, the plate spectra; see Table 2) scattered in and around NGC 1579, scale is 0B19 pixel1 with a field diameter of 3A2. The central suggesting that a pre–main-sequence (PMS) cluster is associ- field was observed alternately with two flanking fields (15000 ated with the cloud. To check that this sample was not con- east and west) in three 20-point dither patterns of 10 s expo- taminated by dMe stars in the foreground, grism images were suresineachfilter.Flat-fieldimagesfreeofdarkcurrentwere also obtained for two clear fields outside the boundary of generated by subtracting on and off exposures of an incan- L1482, centered at 114s, 29300 and 59s, 91800 with respect descent continuum lamp. The science frames were divided by No. 3, 2004 NGC 1579 1237

Fig. 3.—A 300 s I-band image of NGC 1579 with the H-emission stars identified by their Table 1 numbers. LkH 101 is star 194. these median-combined flat-field images to normalize quantum fields represent 100 s of integration time each. The central efficiency variations. field limiting magnitude is estimated at K 18. The 2MASS2 Sky frames were created by masking out stars in the two data when converted to the MKO system agree well with the flanking fields and then combining them with a median filter. present JHK photometry. This median sky frame was then subtracted from all frames to remove the atmospheric emission contribution. The position 2.3. Low-Resolution Spectroscopy centroids of several bright stars in each frame were measured Classification spectrograms of 41 stars in NGC 1579 were and used to calculate shift offsets between the dither pointings. obtained in early 2003 March with the Multi-Object Spectro- With this information, all 40 images were aligned and stacked graph (GMOS) at the 8 m Gemini North Telescope3 on Mauna with a median filter. The result is a composite image covering Kea. The spectrograph was configured with the R831 grating approximately 8A4 in the east-west direction and 40 in the and OG530 order-blocking filter, providing spectral coverage north-south direction. from roughly 5500–8000 8 (dependent upon the slitlet loca- Comparisons of the fluxes in the same stars in different tion in the focal plane mask) at a resolution of 3000. Basic frames throughout the night demonstrated that the data are of reduction procedures including gain correction, bias subtrac- photometric quality. The aperture photometry was performed tion, flat-fielding, scattered light and sky line removal, and using the DAOPHOT package. Observations of the UKIRT faint standard star FS 116 (Hawarden et al. 2001) were used to 2 set the photometric zero points in the new Mauna Kea filter The Two Micron All Sky Survey (2MASS) is a joint project of the Uni- versity of Massachusetts and the Infrared Processing and Analysis Center system (Tokunaga et al. 2002). In addition to the optical in- (IPAC)/California Institute of Technology, funded by the National Aeronautics formation discussed in x 2.1, Table 1 also contains the near- and Space Administration (NASA) and the National Science Foundation (NSF). infrared magnitudes and colors of all the stars detected above 3 The Gemini Observatory is operated by the Association of Universities for the 3 level in at least two bandpasses. The image stacking Research in Astronomy, Inc., under a cooperative agreement with the NSF on employed in the data reduction results in two sensitivity limits behalf of the Gemini partnership: the NSF, the Particle Physics and Astronomy 0 0 Research Council (UK), the National Research Council (Canada), CONICYT for the field, depending on the location. The central 4 ; 4 is (Chile), the Australian Research Council (Australia), CNPq (Brazil), and represented by a 300 s composite image, while the flanking CONICET (Argentina). 1238 HERBIG, ANDREWS, & DAHM Vol. 128

Fig. 4.—Sample of GMOS spectra of H emitters. Each spectrum is identified by its Table 1 number. Prominent emission lines are labeled, along with Li i k6707 absorption. Also indicated are some of the primary features utilized for spectral classification. spectral extraction were performed using the Gemini IRAF (1998) and measurements of spectral indices as defined by package. Wavelength calibration was provided by CuAr lamp Kirkpatrick et al. (1991). Since all of the target stars are exposures. of types K or M, the primary spectral features used were the Spectral types were determined for 38 of the stars by ref- TiO bandheads, the CaH feature near 6975 8,andtheNai D erence to the atlases of Allen & Strom (1995) and Pickles lines. The spectral types are included in Table 1, with notes

Fig. 5.—Color-magnitude diagram: a V0 vs. (V I)0 plot for NGC 1579 stars. Blue points denote stars of known spectral type, which could be corrected individually for extinction. Red points represent stars of unknown type, which were corrected for the mean extinction of AV ¼ 3:5mag.H emission stars are marked by crosses. The solid line is the Pleiades main-sequence ridge line, the dotted extension the ZAMS of Balona & Shobbrook (1984). The isochrones are from D’Antona & Mazzitelli (1997). Several bright stars are labeled with their Table 1 numbers. Since LkH 101 was saturated on all our I-band images, the I magnitude of Barsony et al. (1990) was used. That one point was not corrected for extinction; see text. No. 3, 2004 NGC 1579 1239

Fig. 6.—A V0 vs. (B V )0 diagram for NGC 1579 for which B-band magnitudes are available. The symbols are the same as in Fig. 5. Several of the stars marked by their Table 1 numbers are discussed in the text. As in Fig. 5, the LkH 101 point has not been corrected for extinction. indicating the presence of the Li i k6707 absorption line. The priate distance modulus of NGC 1579. From the color- H equivalent widths determined from the Gemini data are magnitude diagram, it is apparent that the T Tauri star (TTS) included in Table 2 and are sometimes seen to be significantly population in the cluster (distinguished by the presence of H different from those measured from the grism data. A sample emission) has a substantial age spread, with a median value of these spectra is shown in Figure 4. near 0.5 Myr. Since the reddening vector is nearly parallel to the iso- 2.4. Color-Maggnitude Diaggrams chrones in the V0,(V I )0 diagram (Fig. 5), individual red- In the V0,(V I )0 color-magnitude diagram of Figure 5, dening errors will not greatly affect the inferred age but would stars of known spectral type have been corrected individually impact any mass estimates. The coordinates of the point rep- for extinction, normal main-sequence colors being assumed. resenting LkH 101 in Figure 5 (15.71, 3.65) are the observed All others were corrected by the mean cluster extinction values. Given the standard extinction law, that point moves value (AV ¼ 3:5 mag) and the interstellar extinction relations along a reddening vector to cross the Balona & Shobbrook tabulated by Herbig (1998), where AV ¼ 3:08, E(BV ) ¼ (1984) ZAMS at MV ¼3:1, with AV ¼ 9:5 mag. A reason- 2:43E(V I ). The H emission stars are denoted by crosses. able question would be whether the observed colors are sig- Stars of known spectral types are blue, while stars of unknown nificantly contaminated by the strong emission lines in the V type are red. The solid dark line represents the Pleiades main- and I passbands. The line contributions (about 0.07 mag to V, sequence ridge line, obtained from the photometry of Stauffer 0.19 mag to I ) are estimated in the next section and, when (1984), converted to the Cousins photometric system (Bessell applied, lead to MV ¼2:7, AV ¼ 9:2. Such fine-tuning is & Weis 1987), dereddened, and placed at the distance derived hardly justified considering the other uncertainties in the cal- here for NGC 1579 (700 pc, m M ¼ 9:2). A reddening culation, but one can say that the optical photometry is con- vector indicates the shift corresponding to 1 mag of additional sistent with the presence of an early B-type dwarf. visual extinction. Although all the stars above the ZAMS line The brightest stars in L1482 are absent from Figure 5 for in Figure 5 are PMS candidates, this population is undoubtedly lack of I magnitudes, but UBV data are available (see x 5), and somewhat contaminated by unresolved binaries. so those points appear along the top of the main-sequence line The dashed green lines overlaid on the optical color- in the V0,(B V )0 diagram of Figure 6. Several other stars are magnitude diagram in Figure 5 are theoretical PMS isochrones identified in Figure 6 with their Table 1 numbers. Star 40, computed by D’Antona & Mazzitelli (1997).4 The isochrones which lies at the edge of the field, near the southwest corner of were converted from the calculated (log Te,logL)valuesto Figure 3, is probably background, while stars 49 and 162 are the observational color and magnitude coordinates by fitting to likely foreground. Star D is discussed in x 5. the main-sequence (MS) colors and bolometric corrections A(J H, H K ) color-color diagram of the NGC 1579 tabulated by Kenyon & Hartmann (1995), using the appro- cluster is shown in Figure 7, with the symbols having the same meaning as in Figure 5. The solid curves connect normal MS colors as tabulated by Tokunaga (2001) but converted from the 4 There are significant discrepancies between the theoretical PMS iso- Johnson-Glass photometric system to the new Mauna Kea chrone models of different groups. This particular model was chosen only as system using transformations given by Leggett et al. (1992), an example. Carpenter (2001), and Hawarden et al. (2001). The dashed 1240 HERBIG, ANDREWS, & DAHM Vol. 128

Fig. 7.—A J H vs. H K color-color diagram for stars in NGC 1579. The symbols are the same as in Fig. 5. Green points correspond to bright stars in Table 4 (including LkH 101) whose photometry was taken from 2MASS, but converted to the MKO system by the transformations of Carpenter (2001) and Hawarden et al. (2001). The solid curves indicate colors of normal dwarfs and giants, while the dashed lines enclose the normal reddening band. lines define the boundaries of the reddening band within which Whittet (1990) was used instead, the slope would become 1.64, stars with normal (both luminosity class V and III) colors thereby decreasing the number of K-excess stars by 26%. would lie for any extinction. The slope of these lines is defined Kenyon & Hartmann (1995) showed that comparisons be- by the interstellar extinction law of Whittet (1988) for diffuse tween theoretical PMS isochrones and near-infrared color- clouds. A reddening vector is also plotted to mark the shift for magnitude diagrams are not reliable means to derive a cluster 5 visual magnitudes of extinction. The stars that lie above the age because in such a diagram the reddening vector lies nearly reddening band all have photometric errors that could place perpendicular to the isochrones. Therefore, without knowing them within the band. There are 58 stars that lie to the right the amount of extinction each individual star suffers, it is not of the reddening band. Such an excess at 2.2 misregarded possible to assign ages. Further complicating any such analysis as evidence of a circumstellar disk, based upon spectral energy are the large variations in extinction over small angular scales distribution models of stars with known disks. A comparison in NGC 1579. It is worth noting that eight of the (H K )- with W(H) has demonstrated that an H K excess (i.e., the excess stars have J K > 4,and18have3 J K 4(of abscissa distance from the rightmost reddening band border) which four of the former and six of the latter had already been can also discriminate between weak-line T Tauri stars noted by Aspin & Barsony 1994, although the area of their (WTTSs: H emission mostly from chromospheric activity) survey and of ours only partly overlap). Such a concentration andclassicalTTauristars(CTTSs:H emission mostly from of heavily reddened stars suggests that an embedded popula- disk accretion processes; Hartmann 1998). However, in this tion, even younger than that represented in Figures 5 and 7, field the vast majority of H emission stars lie within the may exist in L1482, as was first urged by Barsony et al. (1991), reddening band in Figure 7. If K L colors were available for although the precise size and nature of that population remains all these objects, the issue could be better addressed. Aspin & to be explored. Barsony (1994) did measure L magnitudes for 10 of the heavily reddened stars in Table 1 and found that almost all fell outside 3. THE SPECTRUM OF LkH 101 the J K, K L reddening band, but none of those 10 stars Previous spectroscopy of LkH 101 was almost entirely are known to have H emission. limited to wavelengths longward of about 0.7 m on account In similar analyses in the literature, a standard interstellar of the star’s redness. The spectra to be discussed here cover the extinction law (e.g., Cardelli et al. 1989) has been used, al- region from about 0.43 to 0.68 m (with interorder gaps) and though this may not be correct for dense clouds (Chini & were obtained with the HIRES spectrograph at the Keck I Wargau 1998; Martin & Whittet 1990). Extinction laws de- telescope5 on three occasions: 2000 February 2, November 5, rived for star-forming regions often have RV 4:0ratherthan and 2002 December 16. The FWHM resolution, as measured the conventional 3.1. This results in a shallower slope for the from thorium comparison lines, is 7.0 km s1. Exposure times reddening band in near-infrared color-color diagrams and thus rangedupto40minutes. can exclude stars which might otherwise have been interpreted as showing a disk signature. For instance, the slope of the 5 The W. M. Keck Observatory is operated as a scientific partnership among reddening band in Figure 6 is 1.74, corresponding to normal the California Institute of Technology, the University of California, and NASA. interstellar extinction (Whittet 1988). However, if the extinc- The Observatory was made possible by the generous financial support of the tion law derived for the Oph star-forming region by Martin & W. M. Keck Foundation. No. 3, 2004 NGC 1579 1241

Fig. ii Fig. 8.—HIRES spectrum of LkH 101 (on 2000 February 2) in the 6310– 9.—Fe emission lines kk6247, 6249 in LkH 101 on the three dates 6395 8 region. indicated, showing the variation in line structure near the peaks. The position of the whole line did not change signiﬁcantly.

The spectra were extracted by IRAF routines in a window 200 wide perpendicular to the dispersion, the sky being removed by nent at 6kms1 and the longward at about 13 km s1.The sampling the background on either side of the star, held central relative intensities of these peaks vary; Figure 9 shows how on the 700 long slit. The signal-to-noise ratio per pixel is low on the structure of the Fe ii kk6247, 6249 lines changed between these spectrograms, typically about 20 at 5300 8 and 80 at the three epochs. The position of the line center appears not to 6700 8. vary: the mean Fe ii velocity at the three dates of observation, A continuous spectrum is present, but no absorption lines from about 40 unblended lines, is 4:0 0:2, 3:9 0:2, and are detectable. The Balmer lines of H (discussed in x 3.1) are 4:5 0:3kms1, so the changes in the Fe ii peak structure are prominent in emission. Unfortunately, He i k5875 lies outside probably caused by variations of the relative intensity of the the region covered, and while k6678 is prominent in emission, two overlapping components. it is unfavorably located on the detector. The only other He i The [Fe ii] lines show no such double-peaked structure. line, k5015, is weak, while kk4471, 4921 and He ii k4685 are They appear single and symmetric. On the same three expo- not detected. O i k6046 and Si ii kk6347, 6371 are present. sures as above, the average [Fe ii] velocities were 2:9 0:2, The spectrum of LkH 101 otherwise is dominated by narrow 3:4 0:3, and 3:7 0:4kms1.The[Oi], [N ii], and [S ii] emission lines of Fe ii (about 70 have been measured) plus a lines also appear single, at 2:2 0:3kms1 in the mean. few of Ni ii and Mn ii and [Fe ii] (some 40 have been iden- However, the [O i] lines are not symmetric. They can be tified) and [Ni ii]. No Ti ii lines have been found. A repre- represented by two overlapping Gaussians at 3.5 and 10.5 km sentative sample of those lines are shown in Figure 8, of the s1,eachofFWHM¼ 16 km s1 following correction for 6310–6395 8 region. A list of all the stronger lines measured instrumental resolution. The longward component was about is given in the Appendix. 6% stronger on 2002 December 16, as was the case for the Fe ii Can these emission lines contribute significantly to the lines on that date. The [O i] splitting is probably not resolved as photometry? Their contribution to the V magnitude can be clearly as Fe ii on account of the greater kinetic velocity of estimated by adding up the equivalent widths in the V pass- oxygen. band, weighted by the transmission function of the V filter Quite unlike Fe ii,theSiii lines have extended wings (as is (Straizˇys 1992), and divided by the slope of a Planck function apparent for kk6347, 6371 in Fig. 8), and are centered at 3:5 across the passband at a temperature assumed for the under- 0:8kms1.Theycanbefittedbythesumoftwoalmostco- lying continuum. The result is weakly dependent on the latter; incident Gaussians: in the case of k6347, the central peak by a it is near 0.07 mag. The present spectra do not extend be- narrow component of FWHM ¼ 56 km s1, the wings by a broad yond 6800 8, so the same procedure was followed with the feature of 0.2 that peak intensity and FWHM ¼ 165 km s1. equivalent widths published by Hamann & Persson (1989) and There is an interesting anomaly involving Mn ii.Linesof the transmission function of the I filter (Bessell 1983). This multiplet RMT 13 near 6130 8, shown in Figure 10, appear result is somewhat more sensitive to the continuum slope, but in moderate strength, yet the strongest line of RMT 11 at a reasonable estimate of the emission line contribution to the I 5302.32 8, of comparable upper excitation potential, is not de- magnitude of LkH 101 is about 0.19 mag. This was taken tected. The somewhat weaker lines at 5299.28 and 5296.97 8 of into account in x 2.4. RMT 11 would be masked by Fe ii and [Fe ii] lines near those The Fe ii (and Ni ii and Mn ii) lines are centered at 4.0 km positions, but there is no trace of still weaker RMT 11 lines at s1, but they are double-peaked, with the shortward compo- 5295.29 and 5294.21 8. This same anomaly has been observed 1242 HERBIG, ANDREWS, & DAHM Vol. 128

Fig. 11.—Region of the H line in LkH 101on2002December16.The 8 Fig. 10.—Spectrum of LkH 101 between 6110 and 6140 8, an average of section 6575–6620 has been expanded to show that the longward wing of 8 1 three HIRES spectrograms. The wavelength scale is not corrected for the star’s the H line extends to at least 6600 , or a displacement of 1700 km s . ii ii velocity. Lines of the Mn ii multiplet RMT 13 are indicated. Note that the [N ] k6583 line is significantly narrower than the Fe ; all the forbidden lines in this spectrum are single and narrower than the double permitted lines of Fe ii,Niii,andMnii. Note the diffuse interstellar band at 6613 8.H is shown at larger scale in Fig. 12. in Car by Johansson et al. (1995). Their explanation is that the upper levels of RMT 13 are overpopulated by UV line coinci- dence: multiplet Mn ii UV 15 has the same upper levels as RMT et al. (1976) from the shape of the radio continuum of their 8 11, and its members between 1188 and 1197 overlap a group ‘‘core’’ source; their estimate was n ¼ 1:3 ; 107 cm3. ii 8 e of Si lines (UV 5) between 1190 and 1197 . The optical The FWHM of the continuous spectrum perpendicular to 1 emission lines of Si ii in Car have FWHM near 500 km s , the dispersion ranged from 0B7to1B3 on these HIRES expo- from which Johansson et al. infer that the widths of those UV sures. There was no indication that the stellar emission lines, ii Si lines are sufficient to compensate for the wavelength mis- permitted or forbidden, had any greater extension. ii ii matches with the UV Mn lines. However, the optical Si lines If the continuous spectrum of LkH 101 were that of a B0– of LkH 101 are much narrower, so the UV wavelength mis- B1 star as implied by the radio evidence (x 1), then one would matches would be more serious. The presence of the RMT 13 expect to find high-temperature absorption lines characteristic anomaly in LkH 101 shows that the effect, whatever its ex- of that type. No such lines have been found. In particular, the planation, is operative, although there is one departure from O ii,Niii,andCiii blends between 4638 and 4651 8 and He ii 8 ii expectation: the line at 6130.90 , presumably a blend of Mn k4685 are not detected. They do not coincide with, and so k6130.02 and k6130.92, is approximately twice as strong as could not be concealed by, any low-temperature emission would be expected from laboratory intensities. Possibly there is 6 lines. Any absorption wings on H might have been obliter- another contributor. ated by emission, but at H and H only narrower emission ii 8 The intensity ratio of the [S ] lines at 6716 and 6730 is cores are present, and there is no sign of the broad absorption ii often used as a density indicator in H regions. Both lines are wings of an early B-type MS star. present on two of the HIRES spectrograms. The ratio of equiv- One explanation might be that they are broadened into in- alent widths k6716/k6730 is 0.38 on the 2000 November 5 visibility by electron scattering. Hamann & Persson (1989) spectrogram and 0.36 on 2002 December 16. Those ratios are suggested that the very broad wings on the H emission line 5 3 outside the high-density limit (ne 10 cm )ofthegrids could be explained in that way. They made use of a simple shown by Keenan et al. (1996). The discrepancy is real; it plane-parallel model of Castor et al. (1970) in which a thin cannot be caused by blending of k6730.85 with the nearby layer of free electrons of optical depth e lies above the region [Fe ii] k6729.85 because, although the wings of the two lines where the line spectrum is produced. The observed spectrum overlap, the separation is quite clear at HIRES resolution. seen from above is the sum, weighted by (1 ee )andee , ; 6 3 ii Kelly et al. (1994) derived ne ¼ 1 10 cm from the [S ] respectively. Hamann & Persson found that the best match pair at 1.03 m and similar high values from [Fe ii] ratios. A to their H profile was obtained with an electron temperature high electron density near the star was also inferred by Brown of about 5500 K and e ¼ 0:2. We have followed the same procedure in fitting this model to our own data. The fit depends on the profile assumed for the input emission line, but resulted in almost the same values as Hamann & Persson, 6 It should be mentioned that Mn II RMT 13 is also found in emission in some chemically peculiar B-type stars (Wahlgren & Hubrig 2000). It is not namely 5000 K and 0.15, respectively. apparent how an explanation in terms of non-LTE effects in a stratified Mn With these parameters, as a test, a HIRES spectrogram of layer (Sigut 2001) could apply to LkH 101. the narrow-lined B0 V star BD +46 3474 in the 4620–4680 8 No. 3, 2004 NGC 1579 1243

TABLE 3 LkH 101 H Equivalent Widths

W Date (8) Reference

1970–1971...... >400 Herbig (1971) 1976...... 588 Cohen & Kuhi (1979) 1983...... 1050 Hamann & Persson (1989) 1999–2000...... 464 Hernańdez et al. (2004) 2002 Dec...... 550 50 HIRES, this paper 2003 Jan...... 1075 WFGS, this paper region and near H was processed in the same way. Not surprisingly, in the electron-scattering layer the group of strong absorption lines between 4638 and 4654 8 was reduced to a broad, shallow blend of central depth Ac ¼ 0:024. But at e ¼ 0:15 the fractional contribution of this layer to the combined spectrum is small, only 0.14. If this scattered component were to dominate and thereby account for the nondetection of these absorption lines on our HIRES spectrogram of LkH 101 (which is very noisy in that region), trials show that a e value of 5 or greater would be required. That would be quite incompatible with the fit to H.There- fore, the apparent absence of a B-type absorption spectrum cannot be explained in this way. Fig. 12.—H and H lines of LkH 101 (on 2002 December 16). Note the different intensity scales. The vertical line marks the mean velocity given by 3.1. The Hydroggen Lines the narrow emission lines (3.6 km s1). The H emission line is so strong that it was saturated on the 7 first two spectrograms of this series. On the third night (2002 H,98 for H as compared with the theoretical optically December 16) shorter exposures were obtained. Figure 11 thin Case B relative intensities of 3.0, 1.0, and 0.46 (for 4 shows the spectrum on that date, with the vertical scale ex- Te ¼ 10 K). If these Balmer lines in LkH 101 were also panded for the section longward of H. The wing of the H optically thin, the observed decrement can be explained if the emission line extends to at least 1500 km s1,asnotedby emission-line region is more heavily reddened than the con- Hamann & Persson (1989) and discussed above. The two tinuum to which the Ws are referenced. The difference in AV wings appear quite symmetric when reflected about the line between the two regions that would reconcile the two decre- center. ments depends on the energy distribution of the continuum The measured equivalent width (W )ofH depends criti- (for conversion of Ws to fluxes) and on the reddening law. For cally on how the continuum level is interpolated in from the example, if the reddening law were conventional interstellar edges of this very broad emission line. Different trials produce (RV ¼ 3:1) and if the continuum is that of a 30,000 K black- W-values that scatter between 500 and 600 8.Table3isa body, then AV would be about 1.9 mag. But at the other collection of published and our own W values. Their disper- extreme, if the continuum is produced by dust at 1500 K, then sion probably reflects the difficulty of consistently defining the AV would be about 15 mag. Clearly, an AV obtained from H continuum level. line ratios will be larger than that which measures the fore- The structure near the peak of both H and H is shown ground extinction. expanded in Figure 12, both on a velocity scale. The ordi- These considerations would be irrelevant if the Balmer lines nates are flux differences above the continuum, in units of are optically thick, in which case the Balmer decrement could the continuum. Note that the vertical scales are not the same. be steeper. The vertical bar at 3.6 km s1 marks the mean emission-line In the radio, the slope of the continuum of the core com- velocity. ponent of LkH 101 indicates that it is optically thick at Both lines are double-peaked, but not at the velocities or at wavelengths k0.3 cm (Brown et al. 1976). From this a radial the separation of the Fe ii lines: the H separation is 54 km s1 dependence of ne / r can be inferred, and in this way values as compared with about 19 km s1 for Fe ii. The H-line du- of lying between 2and2.6 have been obtained (Olnon plicity could be (1) the result of the overlapping of two separate 1975; Wright & Barlow 1975; Panagia & Felli 1975). Such a but distinct components (unlikely because there is nothing else radial density profile could be maintained by the outflow of a in the emission spectrum at those velocities), or (2) a single line spherical ionized wind. However, there is no evidence at H, that is divided by a near-central reversal produced by higher or elsewhere in the optical region, of the P Cygni structure that and cooler H i. That, however, does not explain why the central would be expected from such outflow models. No optical minima are at different velocities (11 km s1 for H,2km counterpart has been found of the absorption component at s1 for H) or why the relative peak heights are not the same, 400 km s1 in Br reported by Thompson et al. (1976). Nor or why H isasingleemissionlineat0kms1. The H lines present another interesting anomaly. If their 7 Cohen & Kuhi (1979) measured the same lines with a spectrum scanner as strengths are expressed as equivalent widths, the Balmer 587.5, 34.7, and 4.8 8. Hernańdez et al. (2004) found 464.1, 37.7, and 9.0 8, decrement is very steep: W ¼ 550 50 8 for H,398 for respectively, and 6.7 8 for H. 1244 HERBIG, ANDREWS, & DAHM Vol. 128 is any P Cygni structure present in Br on a recent high-quality K-band spectrogram at 10 km s1 resolution obtained with the NIRSPEC spectrograph at Keck II and kindly made available to us by T. Simon. The narrowness of the Fe ii and other emission lines is also incompatible with their origin in a high-velocity outflow. Thus, the 2envelopeofLkH 101 must be explained in some other way. 4. THE SPECTRUM OF THE NGC 1579 NEBULA Early assertions were that the bright nebulosity ‘‘shows the H and K lines in absorption’’ as in type F (Herbig 1956), and that the ‘‘only certain absorption features [in the star] are the H and K lines of Ca ii and very weakly the G band,’’ suggesting late F (Allen 1973). Note that both those early observations were photographic and made at a time when allowance for sky contribution was very subjective. To clarify this issue, LkH 101 and NGC 1579 were reobserved in 1999 November with the HARIS slit spectrograph at the UH 2.2 m telescope on Mauna Kea. The detector was a Tektronix 20482 CCD. The spectra cover the region from about 3900 to 6000 8.The nominal dispersion was 90 8 mm1; the actual resolution as measured from a night sky emission line was 9 8.The24000 long slit was placed through LkH 101 and rotated so as to cross the brightest section of NGC 1579 to the northeast. The upper panel of Figure 13 shows the spectrum of LkH 101 itself, the lower panel the summed spectra of a strip of the Fig. 13.—Spectrum of LkH 101 (top) and a bright section of NGC 1579 nebulosity between 1500 and 2800 from the star at P.A. 21,both (bottom), obtained with the HARIS spectrograph. The intensity scales as following sky subtraction. The level of the continuous spec- plotted do not represent the true relative brightnesses of the two sources. trum has been set to 1.0 throughout both spectra. The strong H line is off scale on both plots, so it has been truncated. The much superior spectra no persuasive evidence is seen for the F- vertical scale of the lower (nebula) plot has been expanded so type absorption spectrum claimed by Herbig (1956) or by Allen that H is of the same height in both. The very strong airglow (1973). But neither does one see the interstellar Ca ii absorption lines at 5577 and 5893 8 wereonlypartiallyremovedbythe lines at H and K that one would expect to be present, judging sky subtraction, so they have been blanked out. from the strength of interstellar Na i on the HIRES spectra of To first approximation the spectra of star and nebula are much LkH 101. Those early low-resolution photographic claims the same; whether some minor differences between individual should be dismissed as due to inadequate allowance for sky, or emission lines are real would require a detailed investigation on possiblyashavingbeenmisledbygapsbetweenclumpsof better material. This demonstrates that the spectrum of LkH emission lines. The presently available HIRES spectrograms do 101 as seen directly, and as emitted in another direction and then not extend to the H and K region, and are too noisy at 4300 8 to scattered off the nebula, do not differ significantly. On these pronounce upon the G band.

TABLE 4 Coordinates and Lick UBV Photometry of Bright Stars in L1482a

Star VB VU B SpT v veq sin i

A*...... 4 29 28.6 35 13 19 12.72 1.10 0.55 B9 8 150 F* ...... 4 29 52.5 35 22 24 12.65 1.08 0.56 B7 8 180 B*...... 4 29 57.0 35 10 59 12.70 1.06 0.42 B7 2: 200 E*b...... 4 30 10.9 35 19 23 10.96 0.87 0.10 B4 12 90 D*c ...... 4 30 19.2 35 17 46 13.23 1.84 1.08: K0: 9 80 H*d ...... 4 30 51.0 35 26 45 11.49 0.87 0.22 B8 10: 250 Ce...... 4 30 00.4 35 12 39 14.60 1.12 0.53: ...... G...... 4 30 16.4 35 25 22 12.76 1.04 0.63 ...... 101*f...... 4 30 14.4 35 16 24 15.71 2.20 ......

a These UBV observations were obtained at Lick Observatory by B. Paczynski on 1962 December 26 and 1963 January 2. Those stars that illuminate individual reﬂection nebulae are marked with asterisks. The coordinates are in standard units, equinox J2000.0. The velocities quoted in the last two columns are given in units of kilometers per second. A colon after a number marks a particularly uncertain value. b HD 279899. It is star A of Barsony et al. (1990). The V, B V values in their Table 2 must be in error. It is S222, star 1 of Hunter & Massey (1990); they classiﬁed it as B3–7 V. c This is star B of Barsony et al. (1990), for which they give V =13.92,B V = 1.34. Our values (Table 1/249) are 13.43, 1.87. d The nebulosity at H is IC 2067. e Our values (Table 1/40) are V =14.63,B V =1.06. f The V and B V for LkH 101 are from Table 1. Published values are 15.67, 1.99 from Barsony et al. (1990), and 14.76, 0.90 from Bergner et al. (1995). No. 3, 2004 NGC 1579 1245

5. OTHER STARS IN L1482

There are six other fairly bright stars (V ¼ 11 13) within about 150 of LkH 101 that illuminate small local patches of nebulosity and thus must lie in or very near L1482. They are listed in Table 4. Most are saturated on our photometric exposures, so we rely on the UBV data in the table, which were obtained at Lick Observatory in 1962–1963 by B. Paczynski, to whom we are indebted for this unpublished information. HIRES spectrograms of these stars were obtained in 2002 December and 2003 December, and although underexposed, deserve individual description. The spectral types (but not the luminosities) of stars A, F, E, B, and H can be estimated from the He i k4471/Mg i k4481 ratio on the HIRES spectrograms, that ratio being type-sensitive in mid- and later B’s. The results are in Table 4.8 Star D (=HBC 391) has an interesting spectrum. On a low- resolution HARIS spectrogram of the 4000–4500 8 region the type appears to be early K, but in the red it was classified by Herbig & Bell (1988) as K7(Li) with weak H emission: W(H) ¼ 2:3 8. Emission is marginally detectable on the grism spectrograms of 1990 and 1998, and is thus compatible with a line of about that strength. Becker & White (1988) reproduce a low-resolution spectrum (it is their ‘‘lick3’’) Fig. 14.—Top: The 6540–6620 8 region in star D (on 2002 December 16), showing a late-type absorption spectrum and H in emission showing the unusual structure of H (and the strong DIB at 6613 8). Bottom: 1 from which they inferred that D is a TTS. The remarkable Spectrum of the G9 III–IV star HR 1787, here spun up to veq sin i ¼ 80 km s , nature of the spectrum of D becomes apparent at HIRES to match approximately the line widths of star D. resolution, which shows that star D is not a conventional TTS. Li i k6707 is strong (W ¼ 0:26 8). The absorption lines are 1 If it can be assumed that an average extinction of about broad, corresponding to veq sin i about 80 km s , but many A ¼ 3:5 mag and a distance modulus of 9.2 mag (both are asymmetric. The average line centroid is at a velocity of V 1 1 explained in x 6) apply to star D, then MV ¼þ0:5, not far from 9kms , but the line minima lie about 15 km s farther that of a K-type giant. The spectral similarity is apparent in longward. Figure 14, where the lower section shows the spectrum of HR H has a most unusual structure: see the upper section of 1 1787, type G9 III–IV, spun up to veq sin i ¼ 80 km s .Itis Figure 14. The two emission components peaking at about plotted as the point D in Figure 6. 146 and 138 km s1 appear to be the wings of a single, broad 1 Rotational velocities of the B-type stars were estimated by emission line extending to about 300 km s in both directions matching the profiles of their k4471 and k4481 lines with the from the center. The center of this emission line is missing, artificially spun-up HIRES spectrum of the sharp-lined B8 III either because of the superposition of a broad absorption line star HR 562. The resulting values of veq sin i are in Table 4. which itself has an off-center emission core at about 10 km The stars’ radial velocities (v ) in Table 4 are not determinable s1, or because of two separate absorption lines at 44 and 1 with any precision at the higher veq sin i values; those uncer- 49 km s , the ‘‘emission core’’ simply being the gap between. tainties are several km s1. The latter may be the correct explanation because H is only a 1 The spectrum of star E is not quite normal: lower excitation single asymmetric absorption line centered at about 13 km s , lines are symmetric but lines of He i,Siii,andFeiii are not, possibly a blend of the two features seen at H;noemissionis their longward wings being conspicuously steeper than the present at H. shortward. The two emission components of H have equivalent 8 It will be appreciated that in each case these results rest on a widths of about 0.34 and 0.16 , which explains the near- single HIRES spectrogram, so that some of the line asym- undetectability of H emission on the grism spectrograms, metriesinstarsDandEmightbeduetoanimperfectlyre- especially since at that resolution the emission and central 8 solved second spectrum. absorption would run together, with a total W of only 0.29 . The velocities of the five B-type stars and of star D are near However, rough measurement of the Becker & White plot 1 8 enough to the CO velocity of the L1482 cloud (6 km s )that gives W(H) 6 . Given the unusual nature of the spectrum it is unlikely that they are interlopers from the field. Hence and the fact that the object is variable at 3.6 cm (Stine & they are probably products of the same star-forming activity O’Neal 1998) it would not be surprising if the H strength that has produced LkH 101 and the cloud’s TTS population. varied. According to T. Preibisch (2002, private communica- It is interesting that none of these B stars, despite their pre- tion), it is also an X-ray source. sumed youth, show any of the peculiarities that would identify them as HAeBe stars.

8 Photometric spectral types can be obtained from the reddening-independent quantity Q =(UB)0.72(BV ) which is a function of type in the B s. If the 6. THE MOLECULAR CLOUD AND THE DISTANCE Q-values of Schmidt-Kaler’s luminosity class V are adopted, the resulting pho- OF NGC 1579 tometric types agree with the spectroscopic within 1/10 of a spectral class except for star H, where the photometric type is B6 vs. the spectroscopic B8. The Herbig (1956) originally estimated the distance as about photometric types of stars C and G are B8: and B9, respectively. 800 pc on the basis of UBV data and spectral types of two 1246 HERBIG, ANDREWS, & DAHM Vol. 128

TABLE 5 Radial Velocitiesa

v Feature Source (km s1) References

Optical emission lines...... 1 .2 to 4 This paper Br...... 2 .5 2 Simon & Cassar (1984) Br ...... 2 .6 2 12CO (4.7 m) absorption ...... 10 6.9 0.5 Mitchell et al. (1990) H emission ...... 9 2.9 0.7 Fich et al. (1990) Peak of H92 emission...... 3 10.3 1.2 Knapp et al. (1976) 13CO (1–0) emission ...... 4 . 7:: 12CO (1–0) absorption...... 4 . 7: Hi(21 cm) emission...... 5 . 8.7 Dewdney & Roger (1982) Hi(21 cm) absorption ...... 6 . 7 Dewdney & Roger (1986) 12CO (1–0) emission ...... 7 . 6.0 Redman et al. (1986) HCN emission...... 8 . 5.5 Pirogov (1999)

Note.—Sources: (1) The HIRES spectra were extracted with a window 200 wide centered on LkH 101. (2) Aperture diameter 3B75 centered on the star. (3) Emission summed over an H ii source 20 in diameter centered on the star. (4) The 12CO and 13CO emission is from a region 20 –30 in diameter over which line structure is present between 3and11kms1. The narrow 12CO absorption at 7 km s1 is seen against that background. (5) The velocity is an average over an extended region about 50 in diameter. (6) The narrow absorption line is seen against both the broad 21 cm emission line of the H i region and the H ii continuum. (7) Redman et al. (1986) attribute this CO emission to an extended (10 –20) ‘‘obscuring cloud’’ in front of LkH 101. (8) Average over a 4500 beam centered on the star. (9) H emission averaged over a 20 aperture centered about 10 southwest of star. (10) Cold CO absorption against the stellar continuum at 4.7 m. a 1 At the position of LkH 101, vLSR = v7.2 km s .

B-type stars that, because they illuminate small reflection actually B2 1, following Becker & White 1988, then its nebulae, were assumed to be involved in the same cloud as distance would be nearly 1.0 kpc.) LkH 101. The types of stars H (the ‘‘anon’’ of Herbig 1971) If the same calculation is repeated for the five B-type stars and E were estimated to be in the range B3–B5, with E some- with their spectral types from the previous section and values what the earlier, on the basis of low-dispersion Lick photo- for main sequence stars from Schmidt-Kaler (1982), the dis- graphic spectrograms. Both were therefore assumed to be B4 V, tances that follow range from 555 pc (star H) to 800 pc (F), for which MV ¼1:5and(BV )0 ¼0:18 from early com- with a straight mean of 691 55 pc. Therefore, a distance of pilations by Blaauw and by Johnson (in Strand 1963), so a about 700 pc is reasonable. distance of about 800 pc followed if RV ¼ 3:0. If this calcula- The line of sight to NGC 1579 passes north of the Tau-Aur tion is repeated with modern values of the same quantities, the clouds, at about 140 pc, and through the Per OB2 association, distance becomes about 700 pc. (If the type of star E were at roughly 300 pc. Obviously, these distances conflict with the

TABLE 6 Interstellar Absorption-Line Propertiesa

Stars

Line A F B E D H LkH 101 HD 183143

W(5889)...... 466 420 463 468 482b 424 406c 682 W(5895)...... 454 396 449 461 460b 389 403c 609 v (5889)...... 9 9 9d 12 13b 11 10 . ... v (5895)...... 9 9 8d 10 11b 10 9 . ... v (DIBs) ...... 8 8 7 8 79. 8 . ...

DIB Equivalent Widths

5849...... 38 46 46 39 18:. 41 46 82 6195...... 42 40 31 42 34 80 6379...... 50 68 56 47 59 70 56 123 6613...... 146 192 174 153 179 216 105 358 6699...... 15 20 20 17 15:. 23 15:. 60 AV ...... 3.6 3.7 3.7 3.3 ? 3.0 9–18 3.9

a Velocities are heliocentric, in kilometers per second; the equivalent widths are in milliangstroms. b Two minima in the Na i lines are at 7 and 16 km s1. The equivalent widths are with respect to the background of the stellar absorption line. c The widths are with respect to the broad underlying stellar emission lines. d Two minima in the Na i lines are at 5 and 13 km s1. No. 3, 2004 NGC 1579 1247

700 pc just inferred. Despite this overlap in the line of sight, TABLE 7 CO velocities (Ungerechts & Thaddeus 1987) show that three Radio Source Correspondences distinct entities are present. The LSR velocity of the Tau-Aur a CO in that general direction is about 6 km s1, while that of the BW88 SO98 Table 1 Notes dense Per OB2 cloud that contains IC 348 and NGC 1333 81... 5 1 ranges from 6 to 10 km s . NGC 1579 lies on the edge of a ... 2 ... 5 broad band of CO that extends from about 4h40m,35,where 93664 it is confused with Per OB2 material, for about 10 to the 54... 2 northwest along which the velocity changes from about 1 ... 5 ... 1 (near NGC 1579) to 7kms1. The distinction in velocity ... 6 ... 1 between this ‘‘northwest feature’’ and the Tau-Aur, Per OB2 ... 7 194 LkH 101 CO is strikingly shown in Ungerecht & Thaddeus’s Figure 3. ... 8 ... 1 ...... The LSR velocity of the CO (and of other indicators in Table 5) 9 1 1 ... 10 ... 1 at NGC 1579 is 1kms , indicating that NGC 1579 is as- ... 11 ... 1 sociated with the northwest feature, not with Tau-Aur or Per 112... 1 OB2. 3 13 241 Star D There is no independent estimate of the distance of the 214... 1 northwest feature, but some information can be derived from 715... 1 the duplicity of the cores of the interstellar Na i lines in two of 6 16 48, 49 3 the bright stars embedded in L1482. Fits of two overlapping 4 ...... Star E Gaussians to the observed proﬁles produced the velocities 2 ...... 5 given in Table 6. The average LSR velocities of the two Na i a 1 (1) Nothing at radio position in either R or K.(2) components, about 2and6kms , are near the CO ve- Nothing at radio position in R; outside area covered in K.(3) locities expected of both Per OB2 and northwest feature ma- There is a close double star (separation 1B0) at the radio terial in that direction. This would indicate that L1482 and position; the brighter, southeast component has marginal H 8 NGC 1579 are beyond, not in front of Per OB2, thus sup- emission [W(H) 1.7 ]. (4) The star at the position has no grism-discernible H emission. (5) Nothing at radio porting a greater distance. position in K; outside area covered in R.

7. THE RADIO SOURCES But to turn the issue around, note how unsuccessful the radio Becker & White (1988) in a 6 cm survey of the region with 0 observations were in detecting known TTS. Table 2 lists 35 H the VLA detected nine point sources within about 3 of LkH emitters (aside from LkH 101) having W(H) 3 8.Ofthe 101, including LkH 101 itself. Subsequently, Stine & O’Neal 35, 13 have W(H) < 10 8 and thus would be considered (1998) at 3.6 cm with the VLA detected a total of 16 sources WTTS, which are believed to be more likely than CTTS to be in the area, seven being common with the Becker & White list. radio emitters (Chiang et al. 1996). None of those 13 were Stine & O’Neal pointed out that, if they assumed the distance detected with the VLA. of NGC 1579 is 800 pc, then the radio luminosities of these Therefore, we conclude that the clustering of radio sources sources are an order of magnitude higher than those of WTTS around LkH 101 may very well be real, but if so, they must be in the nearby Tau-Aur clouds, at about 140 pc. heavily obscured pre–main-sequence objects of unknown na- The radio observers, in seeking correspondences with op- ture. It would be inadvisable to infer a distance to NGC 1579 tical or infrared sources, have relied on two papers that an- by comparing the radio fluxes of these sources with those of nounced the detection of an infrared cluster centered on ordinary WTTS in Tau-Aur. LkH 101 and gave JHK (and some L) photometry for many stars in the area: Barsony et al. (1991) and Aspin & Barsony 8. INTERSTELLAR FEATURES (1994). We have preferred to use our own catalog (Table 1) of optical and infrared detections. Table 7 summarizes the vari- The strength of the interstellar Na i lines and diffuse inter- ous correspondences. stellar bands (DIBs) in all the NGC 1579 stars observed with Consider first the nature of the objects that were detected in HIRES is striking. In fact, some of the stronger DIBs such as radio. Both Becker & White and Stine & O’Neal detected k4428 and k5780 are obvious on even the low-resolution LkH 101 and star D, but we find that optical or infrared HARIS spectra. The line k6613 appears in Figure 11 of the H point sources are present at only two other radio positions, as region of LkH 101 and in Figure 14 of star D. (In the latter follows. case, it is unusual to see DIBs so clearly against the complex BW6=SO16 coincides with Table 1/48, 49, a close (1B0) background of a late-type star.) Table 6 gives the Wsandve- double, the brighter (southeast) component of which has weak locities for the Na i lines (D1 at 5889 8,D2 at 5895 8)andfor [W(H) 1.7 8] emission. It does not appear in Table 2 all measureable DIBs, and the Ws of a sample of the stronger because we regard grism detections below the 3 8 level as DIBs. unreliable. In this case, the emission is certainly real because Table 5 gives AV s inferred from the B V colors of the the star was already published as HBC 390, with W(H) ¼ B-type stars, and WsforthesameDIBsasobservedinthe 2:0 8 and a spectral classification of M0(Li) from a slit reddened B7 Ia supergiant HD 183143. In the diffuse inter- spectrogram reported by Herbig & Bell (1988). Becker & stellar medium, DIB strengths increase in rough proportion to White reproduce a spectrogram (their ‘‘lick6’’) showing a late- color excess. Given the usual dispersion in that relationship, type absorption spectrum with H clearly in emission. the DIBs in these five B stars are approximately as strong as BW9=SO3 coincides with Table 1/66, a star of V = 19.6 one might expect, judging from the ratio of their AV values to without any grism-detectable H emission. that of HD 183143. If an average foreground AV of about 1248 HERBIG, ANDREWS, & DAHM Vol. 128

3.5 mag is applicable to all the brighter stars in L1482, then the additional V extinction of LkH 101, about 6–7 mag, must be local. If essentially all the interstellar features in LkH 101 are produced in the foreground of L1482, then there is no evidence of any additional interstellar contribution by material in the vicinity of LkH 101, despite its large local extinction. The explanation may be that the material very near LkH 101 is depleted in DIB carriers, as has been observed for TTS (Meyer & Ulrich 1984). 9. FINE STRUCTURE OF LkH 101 As already noted, Danen et al. (1995) were able to set an upper limit of 0B34 for the FWHM of the core component of LkH 101 at 10 m. Subsequently, Tuthill et al. (2002) resolved the star in the H and K bands at a resolution of about 0B02 using the Keck interferometer. Their images show it as a Fig. 15.—Predicted profile of a narrow emission line formed in a flat ro- B tating annulus (of the dimensions of the horseshoe-shaped structure observed horseshoe-shaped partial ring about 0 05 (35 AU) across, open by Tuthill et al. 2002). As explained in the text, this particular calculation on the northeast toward a fainter, bluer companion at a dis- assumes that the annulus extends from 12 to 22 AU, is in Keplerian rotation B tance of 0 18. They regarded this as providing ‘‘definitive about a 15 M star, and that its normal is inclined 30 to the line of sight. If the confirmation of the scenario of an accretion disk with a central double Fe ii lines observed in the optical region originate in such a structure, 1 optically thin cavity’’ (Tuthill et al. 2001), the disk being this is their expected profile, following blurring by the 7 km s instrumental resolution. The ordinate scale has no significance. tipped away from the line of sight by P35. Another interpretation of this structure could be that it represents a tipped annulus either in rotation or expansion Doppler velocity vp ¼ vc sin i, and the velocity profile will around the central star. Assume that the Fe ii emission lines in 2 2 1=2 have the form (vp v ) . Figure 15 shows the result of a the optical region originate in such a flat annulus of azi- sample calculation: the annulus was subdivided into 10 sub- muthally uniform surface brightness, and that their duplicity annuli between r ¼ 12 and 22 AU, the inclination was 30, is produced by its rotation. If in Keplerian rotation about a and the individual line profiles were summed and convolved mass M, the circular velocity at r ¼ 0B025 ¼ 17:5AU(at with the instrumental Gaussian. Given the assumptions in- d ¼ 700 pc) is volved, Figure 15 is a fair representation of the observed Fe ii ii and Ni emission line profiles (Fig. 8), but not of those of 1=2 1=2 [O i]andSiii. M 700 pc 1 vc ¼ 27:6 km s : ð1Þ On the other hand, the horseshoe might simply be a per- 15 M d manent zone through which ejected material flows, although then one might expect P Cygni structure on the emission lines, If the horseshoe structure is really such an annulus with its which is not observed. If instead it were an expanding ring normal inclined i to the line of sight, then vc will project as a (which would produce line doubling very much like Fig. 15),

Fig. 16.—Left: Central region of NGC 1579, centered on LkH 101, in the K band. Right: Same region in R. The area shown is about 2A75 on a side, with north up and east to the left. No. 3, 2004 NGC 1579 1249

dense foreground cloud, but the surface of NGC 1579 (here depicted for simplicity as a relatively thin slab) is illuminated directly by the star, although from Earth that surface is visible only outside the shadow of the foreground cloud. That fraction of the light of LkH 101 not scattered off the slab penetrates it and illuminates the ‘‘bar’’ and ‘‘arc’’ structures. But to be detected at Earth, that radiation must reemerge from the slab and so is attentuated twice. It is suggested that this double passage through the slab, coupled with the difference in the scattering cross-section of dust between R and K wavelengths, can at least qualitatively explain Figure 16. Becker & White (1988) reproduced a ‘‘radiograph’’ of the immediate region of LkH 101 constructed from their 6 cm VLA observations (their Fig. 3). It shows a bright band with considerable structure that curls around the star from northeast to northwest, at a separation to the north of about 1000. It is not the ‘‘bar’’ north of the star seen in Figure 16, which is only slightly curved, and is about 1700 from the star. There is some infrared structure still nearer the star, but better K-band material than ours would be required to resolve it from the scattered light of LkH 101.

11. SUMMARY LkH 101, which illuminates the reflection nebula NGC 1579, has an unusual emission-line spectrum. There is cer- Fig. 17.—Schematic arrangement of star and nearby clouds intended to tainly a high-luminosity star in the core of the radio source explain the difference (seen in Fig. 16) between the appearance of the LkH 101 at that position, whose mass is estimated to be about 15 M region in the R and K bands. It is proposed that the star is heavily screened from from its Lyc flux (inferred from the radio continuum spec- the observer by the dark lane that crosses NGC 1579 from northeast to south- trum), from its position in a (V, V I ) color-magnitude west (see Figs. 1 and 3), but that the unscreened star illuminates more distant dust to produce the bright (at shorter wavelengths) surface of NGC 1579. A diagram, and because such a mass is compatible with a dy- fraction of the star’s light passes through that dusty slab to illuminate the ‘‘bar,’’ namical interpretation of the duplicity of the star’s Fe ii and another fraction is scattered back through the slab toward the observer. But emission lines. The faint H-emission stars found in the since extinction in the slab is wavelength dependent, the ‘‘bar’’ will be seen surrounding molecular cloud (L1482) suggest an age of preferentially at longer wavelengths. The weight of the various arrows suggests about 0.5 Myr for that population. Theory suggests that stars the relative amounts of energy passing along those paths. k of mass 10 M will complete their pre–main-sequence evolution while still heavily obscured, because H burning at a projected expansion velocity of 10 km s1, its radius in the interior will begin, and the star will arrive on the would double in P9 yr. This possibility could be checked ZAMS, while heavy accretion is still underway. Therefore, by repeating the Tuthill et al. observation some years in the not until the foreground material is cleared away will such future, but if the structure does change so significantly on so a star become optically detectable, by which time any short a timescale it ought to have had some effect on the T Tauri–like activity should have subsided. We speculate integrated brightness of LkH 101. We are not aware of any that perhaps there is a time in the early evolution of such a evidence of long-term variability.9 massive ZAMS star when some signature of that recent It is concluded that if the metallic emission lines are indeed activity survives, and that in the case of LkH 101—which produced in the structure reported by Tuthill et al., the rotating falls in that mass range—its unusual spectrum may be such annulus explanation appears somewhat the more satisfactory a signature. of the two. Cohen (1980) envisaged LkH 101 as a hot star that has ionized the central volume of a thick, dusty circumstellar shell 10. THE LINE-OF-SIGHT STRUCTURE NEAR LkH 101 fed by ongoing mass outflow, from which ultraviolet photons escape (possibly because of a flattened geometry) to ionize Figure 16 contains two images of the immediate vicinity the outer envelope. This picture was elaborated by others, of LkH 101: left in the K band (2.2 m) and right in notably Simon & Cassar (1984) and Hamann & Persson R (0.7 m). Note that neither the ‘‘bar’’ just north of the (1989). Our spectroscopy extends those results and provides star or the more distant curved ‘‘arc’’ are seen at R. Figure 17 new details. Clearly, the emission spectrum of LkH 101, is a sketch of an arrangement that can explain this striking which must originate near the star, is not produced in a single difference in the appearance of the nebular structure. It is homogeneous region. That is demonstrated by the variety of proposed that, as seen from Earth, LkH 101 lies behind a species and by their different line widths: the [O i] lines are sharp, Fe ii and [Fe ii] wider, Si ii broader still (FWHM values are given in x 3). This is probably the order in which the 9 LkH 101 is listed in the original Catalog of Suspected Variables (as NSV 1618), apparently on the basis of infrared photometry by Strecker & Ney density and the level of ionization increase with decreasing (1974). But those authors state specifically that the star was not variable over distance from the central star. Some of these systematics have the period of their observations, about 9 months in 1973. been observed in other peculiar stars, so they could indicate 1250 HERBIG, ANDREWS, & DAHM an atmospheric structure common to high-luminosity objects Interestingly, since LkH 101 and all the B-type stars in of this kind. L1482 show the diffuse interstellar band spectrum in com- However, LkH 101 does not conform to expectation in parable strength and roughly as expected for AV 3:5mag, several respects. First, if a very young early B-type star is that additional 6–7 mag of extinction in front of and around present, no sign of its absorption spectrum can be found in the LkH 101 can contribute little more. It is likely that the optical region: the continuum is quite smooth and featureless circumstellar material, like the dust at many T Tauri stars, is except for many strong interstellar features. Second, if the DIB-deficient. material surrounding the core of LkH 101 is supplied by One unusual object in L1482 is the bright star D about 20 continuous outflow from a central star, mass-loss rates can be northeast of LkH 101. It is certainly a member of the cloud inferred from the radio spectrum via the theory of Wright cluster because it illuminates its own small patch of reflection ˙ & Barlow (1975). That theory gives M=v1,wherev1 is the nebulosity at the edge of NGC 1579, and its radial velocity terminal velocity of the outflow. Cohen (1980) assumed that agrees with those of the B stars. The spectrum is that of a v1 is given by the extended wings of the H emission line, K-type giant, MV about 0.5, broad absorption lines (veq sin i ¼ ˙ ; 5 1 1 and so obtained an M of about 3 10 M yr (alterna- 80 km s ), strong Li i k6707, and a very complex H tively, the H wings can be explained by scattering of a core structure. It is also a radio and an X-ray source. Well above the profile in a thin blanket of free electrons: x 3.1). Direct evi- cluster main sequence, it may be on a radiative track toward dence of such mass outflow would be P Cygni structure at the ZAMS at type A. H, but there is no sign of that on the HIRES spectrograms. An absorption component due to cool H i is seen in H and H (Fig. 12), but it is near the systemic velocity. The density 2 We are indebted to the National Science Foundation for profile ne / r inferred from the radio spectrum must have another explanation. partial support of this investigation under grants AST 97- Third, the high level of ionization in the core, if the radio 30934 and AST 02-04021. Dahm acknowledges support continuum is interpreted as free-free emission, implies that during part of this time from the NASA Graduate Student it is an H ii region. But the [O iii] lines at 4959 and 5007 8 Researchers Program. We also appreciate unpublished infor- ordinarily characteristic of an H ii region are not present. mation provided us by Thomas Preibisch and by Ted Simon, One notes that both radio and optical evidence (x 3) indicate and helpful advice from Bo Reipurth. And thanks to Karen 6 3 Teramura for constructing Figure 17. ne k 10 cm , but the critical density for the upper state of those [O iii] transitions is 7 ; 105 cm3 (Osterbrock 1989), so collisional deexcitation would be significant, tending to suppress those lines. But the critical density for the same APPENDIX transitions in [N ii] is still lower, so one expects that kk6548 and 6583 would also not be detectable. Yet there is TABLE OF EMISSION LINES IN LkH 101 a line at the position of the stronger of the two, k6583 (see Table 8 lists the stronger emission lines measured on the Fig. 11). Furthermore, the [O ii] multiplet at 7318–7330 8 HIRES spectrograms of LkH 101. The arrangement is as is present according to Hamann & Persson (1989). There- follows: fore, some better explanation for the absence of [O iii]is Columns (1) and (5).—The intensity-weighted mean wave- required. length of the line, not the wavelength of the intensity peak, The spectrum of the brightest area of NGC 1579 appears and not corrected for the stellar velocity. This mean wavelength identical to that of LkH 101, showing that there is no major difference in the spectrum of the source as seen di- will change as the line structure varies. Many entries are followed by one of these characters: rectly or as radiated in another direction and scattered off nearby dust. It would be worthwhile to observe the nebular d: double peaks; spectrum at much higher resolution to see if the finer details t: triple peaks; are identical. br: the line is broad; After LkH 101, the most luminous stars in the L1482 bl: the line is blended with another; cloud are five of types B4–B9, all of them physically asso- as: the line is asymmetric; ciated with the cloud and members of the cloud cluster on the 1: the line measured on only one spectrogram, usually be- basis of their agreement in radial velocity with cloud CO. cause it falls in an interorder gap on the others; *: a note follows at the end of the table. Notably, none appear to be HAeBe stars. If it is assumed that all five have ZAMS MV values and B V colors, a distance Columns (2) and (6).—W, the equivalent width of the entire of about 700 pc follows. Thus, although the line of sight line,inangstroms.LinesofW less than about 0.05 8 are of passes near the Tau-Aur clouds (at about 140 pc) and through low weight. the Per OB2 association (about 300 pc), NGC 1579/L1482 Columns (3)–(4) and (7)–(8).—The likely identification lies well beyond them both. The AV values of the B stars (with RMT number in parentheses) and its laboratory wave- average 3.5 mag, as do those of the emission-H stars in the length. A plus sign means that there is probably another con- cloud for which spectral types are available, although in the tributor. The [Fe ii]andFeii wavelengths are from Johansson latter case there is large dispersion. (1977) and (1978), respectively. If that average AV of 3.5 mag from front and foreground of Some lines of interest such as He i k5875donot L1482 also applies to LkH 101, its additional AV of 6– appear because they fall in interorder gaps on all three 7 mag must originate in the immediate vicinity of the star. spectrograms. TABLE 8 Emission Lines in LkH 101

Mean W Mean W Mean k (8) Ion Laboratory k Mean k (8) Ion Laboratory k (1) (2) (3) (4) (5) (6) (7) (8)

4340.47...... 9.0 H 40.468 5660.31...... 0.09: ? ... 4351.84...... 0.77 Fe ii (27) 51.764 5673.24...... 0.17 [Fe ii] 73.211 4359.35 1 ...... 0.9 [Fe ii] (7F) 59.333 5675.27 bl ...... 0.05 ? ... 4413.80...... 0.54: [Fe ii] (7F) 13.782 5676.37 bl ...... 0.08 ? ... 4416.41: ...... 0.78: [Fe ii] (6F) 16.266 5718.20...... 0.04 [Fe ii] (39F) 18.216 4452.05...... 0.65 [Fe ii] (7F) + 52.098 5746.92 * ...... 0.53 [Fe ii] (34F) + 46.966 4457.98 1 ...... 0.37 [Fe ii] (6F) 57.945 5754.63...... 0.34 [N ii] 54.59 4474.88: 1 ...... 0.17: [Fe ii] (7F) 74.904 5776.81...... 0.05 ? ... 4489.17: ...... 0.46 [Fe ii] (6F) + 88.749 5791.88...... 0.12 ? ... 4491.47...... 0.41: Fe ii (37) 91.407 5795.31: ...... 0.08 ? ... 4508.32 d: ...... 0.55 Fe ii (38) 08.283 5835.45...... 0.23 [Fe ii] 35.449 4515.38 d ...... 0.98 Fe ii (37) 15.337 5842.40: 1 ...... 0.05 ? ... 4520.27 d ...... 1.13 Fe ii (37) 20.225 5870.04 1 ...... 0.13 [Fe ii] 70.020 4522.67 d ...... 0.92 Fe ii (37) 22.634 5885.11 1 ...... 0.18 ? ... 4549.51 d: ...... 1.45 Fe ii (38) 49.467 5893.46: ...... 0.07: ? ... 4576.43...... 0.34 [Fe ii] + 76.393 5901.28: ...... 0.04: [Fe ii] (34F) 01.263 4582.88...... 0.33: Fe ii (37) 82.835 5902.84...... 0.13 Fe ii 02.825 4583.89 d ...... 2.29 Fe ii (38) 83.829 5913.87...... 0.14 [Fe ii] ? 13.258 4629.39 t ...... 2.40 Fe ii (38) 29.336 5923.74: ...... 0.09 ? ... 4639.70: ...... 0.32 [Fe ii] (4F) 39.667 5948.40...... 0.08 Fe ii 48.419 4666.84...... 0.41 Fe ii (37) 66.750 5955.75: bl...... 0.28 Fe ii 55.700 4728.09...... 0.40: [Fe ii] (4F) 28.068 5957.71...... 1.22 Si ii (4) 57.561 4731.47...... 0.31 Fe ii (43) 31.439 5978.92 1 ...... 2.0 Si ii (4) 78.929 4774.74: ...... 0.16 [Fe ii] (20F) 74.718 5991.48 d ...... 0.24 Fe ii (46) 91.368 4814.59...... 0.79 [Fe ii] (20F) 14.534 6001.37...... 0.06 ? ... 4824.15...... 0.22 Cr ii (30) 24.13 6021.02: ...... 0.03: ? ... 4861.38 d ...... 39.0 H 61.332 6040.56 d ...... 0.24 ? ... 4874.52...... 0.23 [Fe ii] (20F) 74.485 6044.14...... 0.07 [Fe ii] 44.076 4889.65...... 0.77 [Fe ii] (4F) 89.616 6046.35...... 0.82 O i (22) 46.44 4898.64...... 0.22 [Fe ii] 98.607 6052.63...... 0.05 ? ... 4905.38...... 0.58 [Fe ii] (20F) 05.339 6061.12...... 0.06 Fe ii 60.991 4923.97...... 2.16 Fe ii (42) 23.930 6103.56...... 0.04 Fe ii (200) ? 03.54 4947.43...... 0.10 [Fe ii] (20F) 47.373 6113.58...... 0.06 Fe ii (46) 13.330 4950.80...... 0.13 [Fe ii] (20F) 50.744 6122.55 d ...... 0.16 Mn ii (13) 44.438 4973.42...... 0.20 [Fe ii] (20F) 73.388 6124.96 bl ...... 0.43 Ni ii 24.910 5005.52...... 0.13 [Fe ii] (20F) 05.512 6126.10: bl...... 0.20: Mn ii (13) 25.855 5015.64: ...... 0.19 He i 15.675 6129.02...... 0.09 Mn ii (13) 28.725 5018.48...... 2.09 Fe ii (42) 18.450 6130.94: bl...... 0.05: Mn ii (13) 30.796 5020.26...... 0.15 [Fe ii] (20F) 20.233 6131.90 bl ...... 0.12 Mn ii (13) 31.917 5022.59...... 0.12 ? ... 6147.82 d ...... 0.29 Fe ii (74) 47.767 5026.79: ...... 0.09: ? ... 6149.30 d ...... 0.28 Fe ii (74) 49.246 5030.62...... 0.14 ? ... 6156.77...... 0.06 O i (10) ? 56.766 5041.01...... 0.49: Si ii (5) 41.026 6158.35...... 0.28 O i (10) ? 58.184 5056.08 br...... 0.89 Si ii (5) 55.981 6161.07...... 0.12 ? ... 5089.26...... 0.18 Fe ii 89.220 6166.00: ...... 0.03 ? ... 5093.63: ...... 0.09 ? ... 6172.71...... 0.23 ? ... 5100.80...... 0.16 Fe ii 00.735 6175.35: ...... 0.06 ? ... 5108.03...... 0.13 [Fe ii] 07.942 6188.53 1 ...... 0.12 [Fe ii] (44F) 88.552 5144.26...... 0.06 ? 6224.70...... 0.05 Fe ii ? 24.640 5145.97: br...... 0.16 O i (39)? 46.096 6231.9 br,bl...... 0.2 ? ... 5149.41...... 0.14 ? ... 6233.57 d ...... 0.58 Fe ii 33.530 5158.06 bl ...... 0.35 Fe ii + 58.074 6238.46 d ...... 0.35 Fe ii (74) 38.386 5158.83...... 1.14 [Fe ii] (19F) 58.777 6239.98: ...... 0.08 Fe ii 39.905 5164.00...... 0.32 [Fe ii] (35F) 63.952 6243.43: ...... 0.23: Ni ii 43.486 5169.06...... 1.76 Fe II (42) 69.000 6247.62 d ...... 0.87 Fe ii (74) 47.545 5177.10: ...... 0.16 ? ... 6248.95 d ...... 1.51 Fe ii (74) 48.889 5178.42...... 0.10 ? ... 6264.36...... 0.08 ? ... 5180.36...... 0.18 Fe ii 80.314 6300.38 1 ...... 4.4 [O i] 00.304 5182.05...... 0.24 [Fe ii] 81.948 6311.84 as...... 0.12 ? ... 5197.63 d ...... 2.03 Fe ii (49) 97.559 6318.02 d ...... 4.9 Fe ii 17.989 5199.16...... 0.22 Fe ii + 99.123 6338.16 d ...... 0.17 ? ... 5200.80...... 0.10 ? ... 6347.14 br...... 3.2 Si ii (2) 47.103 1252 HERBIG, ANDREWS, & DAHM Vol. 128

TABLE 8—Continued

Mean W Mean W Mean k (8) Ion Laboratory k Mean k (8) Ion Laboratory k (1) (2) (3) (4) (5) (6) (7) (8)

5203.62: ...... 0.15 Fe ii 03.643 6353.16...... 0.06 [Fe ii] 53.116 5216.81...... 0.12 ? ... 6357.18...... 0.16 Fe ii 57.165 5220.11...... 0.23 [Fe ii] 20.059 6363.84...... 1.45 [O i] 63.776 5234.68 d ...... 2.00 Fe ii (49) 34.619 6365.15...... 0.19 [Ni ii] 65.104 5237.57...... 0.17 Cr ii (43) 37.34 6369.55: bl...... 0.11: Fe ii (40) 69.464 5247.99...... 0.24 Fe ii 47.952 6371.62 br...... 1.6 Si ii (2) 71.359 5251.30...... 0.14 Fe ii 51.234 6375.77...... 0.13 Fe ii ? 75.791 5254.99 d ...... 0.29 Fe ii 54.928 6383.76 d ...... 3.1 Fe ii 83.721 5261.68...... 0.90 [Fe ii] 61.621 6385.48 d ...... 1.81 Fe ii 85.455 5264.62 * ...... 0.36 Fe ii (48) 64.805 6417.13...... 0.62 Fe ii (74) 16.921 5273.41...... 0.82 [Fe ii] 73.346 6429.09 d ...... 0.20 Ni ii 28.868 5276.06...... 2.50 Fe ii (49) 75.999 6432.76 d ...... 0.26 Fe ii (40) 32.682 5284.19...... 0.34 Fe ii (41) 84.098 6440.47...... 0.08 [Fe ii] 40.400 5291.68...... 0.17 Fe ii 91.666 6443.00 t:...... 1.45 Fe ii 42.951 5296.88...... 0.20 [Fe ii] 96.829 6446.50...... 0.05 ? ... 5299.01...... 0.21 Fe ii 98.844 6448.97...... 0.05: ? ... 5306.16...... 0.14 Fe ii 06.180 6456.40 d* ...... 2.6 Fe ii (74) 56.389 5316.66 d ...... 5.9 Fe ii (48) + 16.624 6465.72...... 0.16 ? ... 5325.61 d ...... 0.36 Fe ii (49) 25.559 6473.91...... 0.08 [Fe ii] 73.862 5333.66...... 0.68 [Fe ii] + 33.646 6482.27: bl,br...... 0.18 [Fe ii] 82.311 5362.93 t:...... 1.14 Fe ii (48) 62.866 6484.17: bl,br...... 0.33 Ni ii 84.083 5376.51...... 0.57 [Fe ii] 76.452 6485.29 bl ...... 0.13 [Fe ii] 85.282 5395.90...... 0.14 ? ... 6491.36 d ...... 1.05 Fe ii 91.250 5402.16...... 0.10 ? ... 6493.11 d ...... 1.58 Fe ii ? 93.034 5405.50: ...... 0.08 ? ... 6495.24: ...... 0.03 ? ... 5412.72...... 0.28 [Fe ii] 12.654 6506.37 t ...... 0.79 ? ... 5414.11...... 0.12 ? ... 6511.19: 1 ...... 0.05 [Fe ii] ? 11.231 5425.31...... 0.54 Fe ii (49) 25.247 6548.18 bl *...... 0.08 [N ii] 48.05 5433.13 1 ...... 0.51 [Fe ii] (18F) ? 33.129 6562 * ...... H 62.817 5445.84...... 0.08 ? ... 6583.50...... 0.29 [N ii] 83.45 5457.78...... 0.16 Fe ii 57.719 6585.35...... 0.05 [Fe ii] ? 84.405 5465.96: ...... 0.11 Fe ii 65.929 6586.78 d ...... 0.89 Fe ii 86.702 5477.30...... 0.14 [Fe ii] (34F) 77.242 6592.19...... 0.07 ? ... 5487.60: ...... 0.11 Fe ii 87.625 6596.53: ...... 0.08 ? ... 5492.51: ...... 0.12: Fe ii ? 92.399 6598.38 d ...... 0.32 ? ... 5495.91...... 0.12 [Fe ii] (17F) 95.824 6602.52...... 0.08 Ni ii 02.461 5507.05...... 0.14 Ni ii ? 07.214? 6618.48...... 0.07 ? ... 5527.42...... 0.60 [Fe ii] (17F) + 27.340 6626.89...... 0.18 Ni ii 26.687 5529.08...... 0.14 Fe ii 29.061 6631.58...... 0.08 Ni ii 31.637 5534.92 d ...... 1.55 Fe ii (55) 34.834 6666.82...... 0.23 [Ni ii] 66.800 5544.09...... 0.10 ? ... 6678.1: * ...... 1–2: He i 78.149 5551.52: ...... 0.12 [Fe ii] (39F) 51,310 6716.50...... 0.04 [S ii] 16.47 5577.61 as...... 0.30 [O i] (3F) + 77.339 6719.72...... 0.07 ? ... 5588.15...... 0.15 [Fe ii] (39F) 88.154 6729.96...... 0.06 [Fe ii] (31F) 29.856 5613.25...... 0.05 [Fe ii] (39F) 13.268 6730.87...... 0.11 [S ii] 30.87 5627.43 as...... 0.08 [Fe ii] 27.249 6746.76...... 0.07 [Fe ii] + 46.529 5655.04 as...... 0.07 [Fe ii] 54.856 6757.22...... 0.03 ? 5658.03...... 0.05 ? ...

Notes.—Line 5264 has a shortward wing, probably Fe ii k5264.176. Line 5746: the measurements include a shortward wing, possibly Fe ii k5746.578. Line 6456: the measurements include a distinct subsidiary peak on the shortward edge at about 6455.69, probably a blend of two weaker Fe ii lines, plus a contribution from O i RMT 9. Line 6548: The [N ii] line is confused by an overlapping atmospheric H2O line. Line 6562: see the text for a discussion of H. Line 6678 falls in a corner of the frame and possibly for that reason measurements are discordant.

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