Lkhα-101 and the Young Cluster in NGC 1579

LkHα-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, Hawaii 96822, U.S.A.

ABSTRACT

The central region of the dark cloud L1482 is illuminated by LkHα-101, a heav- 3 ily reddened (AV ≈ 10 mag) high-luminosity (≥ 8 × 10 L⊙) star having an unusual emission-line spectrum plus a featureless continuum. About 35 much fainter (mostly between R = 16 and >21) Hα 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 5 bright B-type stars in the cloud, presumably of about the same age; none show the peculiarities expected of HAeBe stars. De-reddened, their apparent V magnitudes lead to a distance of about 700 pc. Radio observations suggest that the optical object LkHα-101 is in fact a hot star surrounded by a small H II 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 ZAMS star. The V − I 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 LkHα-101. The Hα emission line is very strong, with wings extending to about ±1700 km s−1, which could be produced by a thin overlying layer of hot electron scatterers. There is no sign of P Cyg-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 s−1. Lines of [Fe II], [O I], [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 0′′. 05 (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 LkHα-101 is a massive star caught in an early evolutionary state.

Subject headings: open clusters and associations: individual (NGC 1579) — stars: individual (LkHα-101) — stars: emission-line — stars: formation — stars: pre-main sequence – 2 –

1. Introduction

Current belief is that the formation of a massive star takes place deep in the parent cloud, behind very heavy extinction, and hence will not be optically accessible (Palla & Stahler 1990; Bernasconi & Maeder 1996). But when such a star evolves further to the zero-age main sequence (ZAMS), its radiation and wind will clear out the neighborhood and the star could at some time become optically detectable. We examine here the possibility that LkHα-101, in the nebula NGC 1579, may be such a transition object.

NGC 1579 is a clump of bright nebulosity about 2′ across, lying in a dark cloud (L1482) north of the main Taurus-Auriga cloud complex, a line of sight that also passes through the more distant Per OB2 association. Although there are several stars illuminating their own small reﬂection nebulae nearby, in the era of photography in blue-violet light there was no obvious source of illumination of NGC 1579 itself, nor was any cluster of embedded stars apparent. In 1956 it was found (Herbig 1956) that at the edge of the bright nebulosity there is a very faint, very red star (V ∼ 16) having a powerful Hα emission line and an unusual emission-line spectrum. Subsequent work has shown that the spectra of star and nebula are identical (§4), and that the polarization of the nebula points to a source at that location (Redman et al. 1986), so this is the illuminating source of the nebula. The spectrum of this star, named LkHα-101, has since then been investigated in increasingly greater detail (Herbig 1971; Allen 1972; Thompson et al. 1976; Simon & Cassar 1984; Hamann & Persson 1989, and others).

Early radio observations (Brown, Broderick, & Knapp 1976; Cohen 1980; Purton et al. 1982) indicated that there were two principal contributors to the radio continuum at NGC 1579. The ﬁrst is a point source at the position of LkHα-101, believed to be a hot star plus a very small H II region, both behind an optically-thick dust shell. Later, Hoare et al. (1994) and Hoare & Garrington (1995), from Merlin interferometry, determined mean sizes for this source of 0′′. 55 at 18.7 cm and 0′′. 31 at 6 cm. At 10 µm, Danen, Gwinn, & Bloemhof (1995) found it unresolved, and set an upper limit of 0′′. 34 on its FWHM. This is the object to which the optical spectroscopy of LkHα − 101 refers. The second radio continuum component is an extended H II envelope, of diameter of the order of 1′, the whole being embedded in a clumpy H I cloud about 5′ in diameter (Dewdney & Roger 1986).

Sharpless (1959) probably included NGC 1579 as S222 in his catalog of H II regions because it is so much brighter on the red Palomar plates than on the blue. However, contrary to that and to expectation from the radio results, the [O III] λλ4957, 5007 lines, characteristic of H II regions, are not present in the spectrum of the nebula or of LkHα-101. Clearly, NGC 1579 is a reﬂection nebula and owes its redness to interstellar extinction and to the nature of its illuminating source. But why that source does not show the spectrum of a H II region is another matter.

If the radio continuum spectrum of the core source in LkHα-101 is the free-free emission of a spherically-symmetric H II region, then the Lyman continuum (Lyc) ﬂux required to maintain that ionization can be obtained (Harris 1976; Brown, Broderick, & Knapp 1976; Knapp et al. 1976). All – 3 – found, on the basis of the properties of OB stars tabulated by Panagia (1973), that the required 1 Lyc ﬂux could be supplied by a single star, if near the ZAMS, of type B0 or B1 . Assuming that the optical continuum is indeed that of a B0.5 ZAMS star, Cohen (1980) found from the narrow-band colors of LkHα-101 that AV = 9.1 ± 0.5 mags. He found also from recombination theory that the Balmer decrement, those lines assumed optically thin, gave AV = 12.5 ± 1 mags. Subsequent investigators (Thompson et al. 1976; McGregor, Persson, & Cohen 1984; Rudy et al. 1991; Kelly, Rieke, & Campbell 1994) using various procedures have found values of AV ranging from 9.7 to 15.8 mags.

If one simply assumes that AV = 10 mag. and a distance of 700 pc (obtained later in this paper) then the observed value of V = 15.7 leads to MV = −3.5. This is not incompatible with a normal B0.5 V, given the crudity of this calculation plus the scatter in the values of MV found in the literature for that spectral type: Panagia (1973) gave −3.5; Vacca, Garmany, & Shull (1996) gave −4.1; while Andersen (1991) found −3.2 and −2.9 from two eclipsing binaries. The mass of a B0.5 V according to the latter two sources is 19 and 13 M⊙, respectively.

In respect to total luminosity, the match with expectation does not seem so satisfactory. The 4 same authorities (above) give for L of a B0.5 V star (in units of 10 L⊙): Panagia 2.0, Vacca et al. 6.2, Andersen 1.9 and 1.4. Actual integration of the narrow-band photometry of LkHα-101 between 1.3 and 25 µm by Strecker & Ney (1974) and by Simon & Cassar (1984), corrected as above for extinction and distance, gives 0.8 in the same units. However, there is warm dust, presumably illuminated by the star, well away from the core. Harvey, Thronson & Gatley (1979) found that 4 between ∼1 and 160 µm this “extended region” has a total L = 1.2 × 10 L⊙, still somewhat low. The explanation may simply be that most of the radiation of an early B star is in the ultraviolet, and such integrations must miss that fraction that is not re-thermalized by circumstellar dust, or does not appear in the near-infrared free-free continuum.

Therefore there is reason to suspect that LkHα-101 may indeed be a star of mass about 15 M⊙ in an interesting phase of its early evolution. We proceed on that assumption.

In what follows, we discuss (§2) our optical and near-infrared photometry of the heavily obscured cluster of stars surrounding LkHα-101, (§3) the optical spectrum of LkHα-101 at high- resolution, and (§5, 6) the surrounding molecular cloud and its contents.

1 More recent calculations (Vacca, Garmany, & Shull 1996) based on later values of Te and L and improved atmospheric models (Sternberg, Hoffmann & Pauldrach 2003) predict substantially higher Lyc fluxes for OB stars, so a somewhat later B type for LkHα-101 would follow. It is not possible to be more specific until such calculations are extended to types later than B0.5. – 4 –

2. The Star Cluster

2.1. Optical Photometry

Figure 1 is a false-color composite of NGC 1579 created from 300 s B, V , and R images of this dataset. Optical images behind BV RcIc ﬁlters (hereafter we omit the subscripts on Rc and ′ ′ Ic), covering an area of about 7.5×7.5 centered on LkHα-101, were obtained at the f/10 focus of the University of Hawaii (UH) 2.2 m telescope in 1999 October. Conditions were photometric, the seeing FWHM about 0′′. 7 (in V ). Exposure times were 10, 60, and 300 s in each ﬁlter so both the bright stars and the faint cluster population were measureable. The detector was a Tektronix 20482 CCD, the scale 0′′. 22 pixel−1.

All the data frames were corrected for bias by subtraction of a median-combined composite dark exposure. Any residual bias was removed from individual frames by fitting a polynomial to the overscan region, and subtracting the fit across the frame. Flat-field images were generated from observations of the twilight sky. Quantum-efficiency variations (typically <1% for this CCD) were removed by dividing the data frames by these median-combined flat-field images.

Aperture photometry was performed using the DAOPHOT package in IRAF. In the rare cases where a star fell on the boundary between bright nebulosity and a darker background, the sky contribution was removed on an individual basis, not the usual annular procedure. Closely grouped stars were analyzed separately using point spread function (PSF) ﬁtting photometry with the PSF + ALLSTAR tasks. The measured instrumental magnitudes were converted to the BV RI system using observations of Landolt (1992) standard stars over a range of airmasses. The limiting magnitude is about V = 22, with completeness to V ≈ 20.5. Table 1 is a list of optical magnitudes and colors for all stars detected above the 3σ level in at least two bandpasses. The coordinates were derived by reference to stars in the HST Guide Star and USNO-A Catalogues. They are reliable at about the 1′′ level. The internal errors of the photometry are shown in Figure 2.

The field was re-imaged again with the same instrumentation on 2003 November 10, under photometric conditions but with only average seeing and brightly moonlit sky. Good photometry was possible for only the brighter stars in the field, but was sufficient to confirm the photometric zero points of the earlier observations.

The heavy and irregular extinction in NGC 1579 causes the optical sample of Table 1 to be an incomplete representation of the faint cluster population. Nearly 100 stars that were detected only in I are not listed there.

Slitless grism spectra (6300 - 6700 A)˚ covering the same area were obtained at the UH 2.2 m telescope in 1990, 1998, and 2003. The detectors, scale, etc. were as described by Herbig & Dahm (2002) in an investigation of IC 5146 with the same equipment. On these exposures about 19 faint stars having Hα in emission were found (16 others were later found with GMOS spectra: see below), scattered in and around NGC 1579, suggesting that a pre-main sequence (PMS) cluster is – 5 – associated with the cloud. To check that this sample was not contaminated by dMe stars in the foreground, grism images were also obtained for two clear fields outside the boundary of L1482, centered at +114s, −293′′ and −59s, −918′′ with respect to LkHα-101. No Hα emitters were found in either of these fields. Equivalent widths of the Hα line in all the Hα emitters (grism + GMOS detections) are given in Table 2. They are assigned IHα numbers in continuation of the numbering system of Herbig & Dahm. In Figure 3 all these stars are identified with their numbers in Table 1.

There are very few Hα-emitters, but many heavily-reddened stars, east of LkHα-101, probably because of greater foreground extinction in that area, as suggested by the CO contours in Barsony et al. (1990, their Figures 11 and 12).

2.2. Near-Infrared Photometry

Near-infrared JHK images of NGC 1579 were obtained in 2002 August with the QUIRC infrared array (Hodapp et al. 1995) at the f/10 focus of the UH 2.2 m telescope. The typical seeing FWHM was ∼ 0′′. 5 (in K). In this configuration, the plate scale is 0′′. 19 pixel−1 with a field diameter of 3.′2. The central field was observed alternately with two flanking fields (150′′ E and W) in three 20-point dither patterns of 10 s exposures in each filter. Flat-field images free of dark current were generated by subtracting on and off exposures of an incandescent continuum lamp. The science frames were divided by these median-combined flat-field images to normalize quantum efficiency variations.

Sky frames were created by masking out stars in the two flanking fields and then combining them with a median filter. This median sky frame was then subtracted from all frames to remove the atmospheric emission contribution. The position centroids of several bright stars in each frame were measured and used to calculate shift offsets between the dither pointings. With this information, all 40 images were aligned and stacked with a median filter. The result is a composite image covering approximately 8.′4 in the E-W direction and 4′ in the N-S direction.

Comparisons of the fluxes in the same stars in different frames throughout the night demonstrated that the data are of photometric quality. The aperture photometry was performed using the DAOPHOT package. Observations of the UKIRT faint standard star FS116 (Hawarden et al. 2001) were used to set the photometric zero points in the new Mauna Kea filter system (Tokunaga, Simons, & Vacca 2002). In addition to the optical information discussed in §2.1, Table 1 also contains the near-infrared magnitudes and colors of all the stars detected above the 3σ level in at least two bandpasses. The image stacking employed in the data reduction results in two sensitivity limits for the field, depending on the location. The central 4′×4′ is represented by a 300 s composite image, while the flanking fields represent 100 s of integration time each. The central field limiting magnitude is estimated at K ≈ 18. The 2MASS2 data when converted to the MKO system agree

2The Two Micron All Sky Survey (2MASS) is a joint project of the University of Massachusetts and the Infrared – 6 – well with the present JHK photometry.

2.3. Low-Resolution Spectroscopy

Classification spectrograms of 41 stars in NGC 1579 were obtained in early 2003 March with the Multi-Object Spectrograph (GMOS) at the 8 m Gemini North telescope3 on Mauna Kea. The spectrograph was configured with the R831 grating and OG530 order-blocking filter, providing spectral coverage from roughly 5500 - 8000 A˚ (dependent upon the slitlet location in the focal plane mask) at a resolution of ∼3000. Basic reduction procedures including gain correction, bias subtraction, flat-fielding, scattered light and sky line removal, and spectral extraction were performed using the Gemini IRAF package. Wavelength calibration was provided by CuAr lamp exposures.

Spectral types were determined for 38 of the stars by reference to the atlases of Allen & Strom (1995) and Pickles (1998) and measurements of spectral indices as defined by Kirkpatrick, Henry, & McCarthy (1991). Since all of the target stars are of types K or M, the primary spectral features used were the TiO bandheads, the CaH feature near 6975 A,˚ and the Na I D lines. The spectral types are included in Table 1, with notes indicating the presence of the Li I λ6707 absorption line. The Hα equivalent widths determined from the Gemini data are included in Table 2, and are sometimes seen to be significantly different from those measured from the grism data. A sample of these spectra is shown in Figure 4.

2.4. Color-Magnitude Diagrams

In the V0, (V − I)0 color-magnitude diagram of Figure 5, stars of known spectral type have been corrected individually for extinction, normal main sequence colors being assumed. All others were corrected by the mean cluster extinction value (AV = 3.5 mags) and the interstellar extinction relations tabulated by Herbig (1998), where AV = 3.08 E(B−V ) = 2.43 E(V −I). The Hα emission stars are denoted by crosses. Stars of known spectral types are blue, while stars of unknown type are red. The solid dark line represents the Pleiades main sequence ridge line, obtained from the photometry of Stauﬀer (1984), converted to the Cousins photometric system (Bessell & Weis 1987), de-reddened, and placed at the distance derived here for NGC 1579 (700 pc, m−M = 9.2). A reddening vector indicates the shift corresponding to 1 mag of additional visual extinction. Although

Processing and Analysis Center (IPAC)/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation. 3The Gemini Observatory is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the Particle Physics and Astronomy Research Council (United Kingdom), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), CNPq (Brazil), and CONICET (Argentina). – 7 – all the stars above the ZAMS line in Figure 5 are PMS candidates, this population is undoubtedly somewhat contaminated by unresolved binaries.

The dashed green lines overlaid on the optical color-magnitude diagram in Figure 5 are theoretical PMS isochrones computed by D’Antona & Mazzitelli (1997)4. The isochrones were converted from the calculated (log Te, log L) values to the observational color and magnitude coordinates by ﬁtting to the main sequence (MS) colors and bolometric corrections tabulated by Kenyon & Hart- mann (1995), using the appropriate distance modulus of NGC 1579. From the color-magnitude diagram, it is apparent that the T Tauri star (TTS) population in the cluster (distinguished by the presence of Hα emission) has a substantial age spread, with a median value near 0.5 Myr.

Since the reddening vector is nearly parallel to the isochrones in the V0, (V − I)0 diagram (Figure 5), individual reddening errors will not greatly aﬀect the inferred age, but would impact any mass estimates. The coordinates of the point representing LkHα-101 in Figure 5 (15.71, 3.65) are the observed values. Given the standard extinction law, that point moves along a reddening vector to cross the Balona & Shobbrook (1984) ZAMS at MV = −3.1, with AV = 9.5 mags. A reasonable question would be whether the observed colors are signiﬁcantly contaminated by the strong emission lines in the V and I passbands. The line contributions (about 0.07 mags to V ,

0.19 mags to I) are estimated in the next section, and when applied, lead to MV = −2.7, AV = 9.2. Such ﬁne-tuning is hardly justiﬁed considering the other uncertainties in the calculation, but one can say that the optical photometry is consistent with the presence of an early B-type dwarf.

The brightest stars in L1482 are absent from Fig. 5 for lack of I magnitudes (as explained in §5) but UBV data are available (Table 4) and so those points appear along the top of the main sequence line in the V0, (B − V )0 diagram of Fig. 6. Several other stars are identiﬁed in Fig. 6 with their Table 1 numbers. Star 40, which lies at the edge of the ﬁeld, near the southwest corner of Fig. 3, is probably background, while stars 49 and 162 are likely foreground. Star D is discussed in §5.

A (J − H), (H − K) color-color diagram of the NGC 1579 cluster is shown in Figure 7, with the symbols having the same meaning as in Fig. 5. The solid curves connect normal MS colors as tabulated by Tokunaga (2001) but converted from the Johnson-Glass photometric system to the new Mauna Kea system using transformations given by Leggett, Smith, & Oswalt (1992), Carpenter (2001), and Hawarden et al. (2001). The dashed lines define the boundaries of the reddening band within which stars with normal (both luminosity class V and III) colors would lie for any extinction. The slope of these lines is defined by the interstellar extinction law of Whittet (1988) for diffuse clouds. A reddening vector is also plotted to mark the shift for 5 visual magnitudes of extinction. The stars that lie above the reddening band all have photometric errors which could place them within the band. There are 58 stars which lie to the right of the reddening band. Such an excess

4There are significant discrepancies between the theoretical PMS isochrone models of different groups. This particular model was chosen only as an example. – 8 – at ∼2.2 µm is regarded as evidence of a circumstellar disk, based upon spectral energy distribution models of stars with known disks. A comparison with W (Hα) has demonstrated that a (H − K) excess (i.e., the abscissa distance from the rightmost reddening band border) can also discriminate between weak-line T Tauri stars (WTTS: Hα emission mostly from chromospheric activity) and classical T Tauri stars (CTTS: Hα emission mostly from disk accretion processes; Hartmann 1998). However, in this field the vast majority of Hα emission stars lie within the reddening band in Fig. 7. If (K − L) colors were available for all these objects the issue could be better addressed. Aspin & Barsony (1994) did measure L magnitudes for 10 of the heavily-reddened stars in Table 1 and found that almost all fell outside the J − K, K − L reddening band, but none of those 10 stars are known to have Hα emission.

In similar analyses in the literature, a standard interstellar extinction law (e.g., Cardelli, Clay- ton, & Mathis 1989) has been used, although this may not be correct for dense clouds (Chini & Wargau 1998; Martin & Whittet 1990). Extinction laws derived for star-forming regions often have RV ≥ 4.0, rather than the conventional 3.1. This results in a shallower slope for the reddening band in near-infrared color-color diagrams, and thus can exclude stars which might otherwise have been interpreted as showing a disk signature. For instance, the slope of the reddening band in Figure 6 is 1.74, corresponding to normal interstellar extinction (Whittet 1988). However, if the extinction law derived for the ρ Oph star-forming region by Martin & Whittet (1990) was used instead, the slope would become 1.64, thereby decreasing the number of K-excess stars by 26%.

Kenyon & Hartmann (1995) showed that comparisons between theoretical PMS isochrones and near-infrared color-magnitude diagrams are not reliable means to derive a cluster age because in such a diagram the reddening vector lies nearly perpendicular to the isochrones. Therefore, without knowing the amount of extinction each individual star suﬀers, it is not possible to assign ages. Further complicating any such analysis are the large variations in extinction over small angular scales in NGC 1579. It is worth noting that 8 of the H − K-excess stars have (J − K) > 4, and 18 have 3 ≤ (J −K) ≤ 4, (of which 4 of the former and 6 of the latter had already been noted by Aspin & Barsony (1994), although the area of their survey and of ours only partly overlap). Such a concentration of heavily reddened stars suggests that an embedded population, even younger than that represented in Figs. 5 and 7, may exist in L1482, as was ﬁrst urged by Barsony, Schombert, & Kis-Halas (1991), although the precise size and nature of that population remains to be explored.

3. The Spectrum of LkHα-101

Previous spectroscopy of LkHα-101 was almost entirely limited to wavelengths longward of about 0.7 µm on account of the star’s redness. The spectra to be discussed here cover the region from about 0.43 to 0.68 µm (with inter-order gaps) and were obtained with the HIRES spectrograph – 9 – at the Keck I telescope5 on three occasions: 2000 February 2 and November 5, and 2002 December 16. The FWHM resolution, as measured from thorium comparison lines, is 7.0 km s−1. Exposure times ranged up to 40 minutes.

The spectra were extracted by IRAF routines in a window 2′′ wide perpendicular to the dispersion, the sky being removed by sampling the background on either side of the star, held central on the 7′′-long slit. The S/N per pixel is low on these spectrograms, typically about 20 at 5300 A˚ and 80 at 6700 A.˚

A continuous spectrum is present, but no absorption lines are detectable. The Balmer lines of H (discussed below) are prominent in emission. Unfortunately He I λ5875 lies outside the region covered, and while λ6678 is prominent in emission, it is unfavorably located on the detector. The only other He I line, λ5015, is weak, while λλ4471, 4921 and He II λ4685 are not detected. O I λ6046 and Si II λλ6347, 6371 are present. The spectrum of LkHα-101 otherwise is dominated by narrow emission lines of Fe II (about 70 have been measured) plus a few of Ni II and Mn II, and [Fe II] (some 40 have been identiﬁed) and [Ni II]. No Ti II lines have been found. A representative sample of those lines are shown in Figure 8, of the 6310 - 6395 A˚ region. A list of all the stronger lines measured is given in the Appendix.

Can these emission lines contribute significantly to the photometry? Their contribution to the V magnitude can be estimated by adding up the equivalent widths in the V passband, weighted by the transmission function of the V filter (Straiˇzys 1982), and divided by the slope of a Planck function across the passband at a temperature assumed for the underlying continuum. The result is weakly dependent on the latter; it is near 0.07 mags. The present spectra do not extend beyond 6800 A,˚ so the same procedure was followed with the equivalent widths published by Hamann & Persson (1989) and the transmission function of the I filter (Bessell 1983). This result is somewhat more sensitive to the continuum slope, but a reasonable estimate of the emission line contribution to the I magnitude of LkHα-101 is about 0.19 mags. This was taken into account in §2.4.

The Fe II (and Ni II and Mn II) lines are centered at +4.0 km s−1, but they are double-peaked, with the shortward component at −6km s−1 and the longward at about +13 km s−1. The relative intensities of these peaks vary; Figure 9 shows how the structure of the Fe II λλ6247, 6249 lines changed between the three epochs. The position of the line center appears not to vary: the mean Fe II velocity at the 3 dates of observation, from about 40 unblended lines, is +4.0 ± 0.2, +3.9 ± 0.2, and +4.5 ± 0.3 km s−1, so the changes in the Fe II peak structure are probably caused by variations of the relative intensity of the two overlapping components.

The [Fe II] lines show no such double-peaked structure. They appear single and symmetric. On the same three exposures as above, the average [Fe II] velocities were +2.9 ± 0.2, +3.4 ± 0.3,

5The W.M. Keck Observatory is operated as a scientiﬁc partnership among the California Institute of Technology, the University of California, and the National Aeronautics and Space Administration. The Observatory was made possible by the generous ﬁnancial support of the W.M. Keck Foundation. – 10 – and +3.7 ± 0.4 km s−1. The [O I], [N II], and [S II] lines also appear single, at +2.2 ± 0.3 km s−1 in the mean.

However the [O I] lines are not symmetric. They can be represented by two overlapping gaussians at −3.5 and +10.5 km s−1, each of FWHM = 16 km s−1 following correction for instrumental resolution. The longward component was about 6% the stronger on 2002 December 16, as was the case for the Fe II lines on that date. The [O I] splitting is probably not resolved as clearly as Fe II on account of the greater kinetic velocity of oxygen.

Quite unlike Fe II, the Si II lines have extended wings (as is apparent for λλ6347, 6371 in Fig. 8), and are centered at +3.5 ± 0.8 km s−1. They can be ﬁtted by the sum of two almost coincident gaussians: in the case of λ6347, the central peak by a narrow component of FWHM = 56 km s−1, the wings by a broad feature of 0.2 that peak intensity and FWHM = 165km s−1.

There is an interesting anomaly involving Mn II. Lines of multiplet RMT 13 near 6130 A,˚ shown in Fig. 10, appear in moderate strength, yet the strongest line of RMT 11 at 5302.32 A,˚ of comparable upper EP, is not detected. The somewhat weaker lines at 5299.28 and 5296.97 A˚ of RMT 11 would be masked by Fe II and [Fe II] lines near those positions, but there is no trace of still weaker RMT 11 lines at 5295.29 and 5294.21 A.˚ This same anomaly has been observed in η Car by Johansson et al. (1995). Their explanation is that the upper levels of RMT 13 are overpopulated by UV line coincidence: multiplet Mn II UV 15 has the same upper levels as RMT 11, and its members between 1188 and 1197 A˚ overlap a group of Si II lines (UV 5) between 1190 and 1197 A.˚ The optical emission lines of Si II in η Car have FWHM near 500 km s−1, from which Johansson et al. infer that the widths of those UV Si II lines are suﬃcient to compensate for the wavelength mismatches with the UV Mn II lines. However, the optical Si II lines of LkHα-101 are much narrower, so the UV wavelength mismatches would be more serious. The presence of the RMT 13 anomaly in LkHα-101 shows that the eﬀect, whatever its explanation, is operative, although there is one departure from expectation: the line at 6130.90 A,˚ presumably a blend of Mn II λ6130.02 and λ6130.92, is approximately twice as strong as would be expected from laboratory intensities. Possibly there is another contributor6.

The intensity ratio of the [S II] lines at 6716 and 6730 A˚ is often used as a density indicator in H II regions. Both lines are present on two of the HIRES spectrograms. The ratio of equivalent widths λ6716/λ6730 is 0.38 on the 2000 November 5 spectrogram and 0.36 on 2002 December 16. 5 −3 Those ratios are outside the high-density limit (ne ∼ 10 cm ) of the grids shown by Keenan et al. (1996). The discrepancy is real; it cannot be caused by blending of λ6730.85 with the nearby [Fe II] λ6729.85 because, although the wings of the two lines overlap, the separation is quite clear 6 −3 at HIRES resolution. Kelly, Rieke, & Campbell (1994) derived ne = 1 × 10 cm from the [S II] pair at 1.03 µm, and similar high values from [Fe II] ratios. A high electron density near the star

6It 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 apparent how an explanation in terms of non-LTE eﬀects in a stratiﬁed Mn layer (Sigut 2001) could apply to LkHα-101. – 11 – was also inferred by Brown, Broderick, & Knapp (1976) from the shape of the radio continuum of 7 −3 their “core” source; their estimate was ne = 1.3 × 10 cm . The FWHM of the continuous spectrum perpendicular to the dispersion ranged from 0′′. 7 to 1′′. 3 on these HIRES exposures. There was no indication that the stellar emission lines, permitted or forbidden, had any greater extension.

If the continuous spectrum of LkHα-101 were that of a B0-B1 star as implied by the radio evidence (§1), then one would expect to ﬁnd high-temperature absorption lines characteristic of that type. No such lines have been found. In particular, the O II, N III, and C III blends between 4638 and 4651 A,˚ and He II λ4685 are not detected. They do not coincide with, and so could not be concealed by, any low-temperature emission lines. Any absorption wings on Hα might have been obliterated by emission, but at Hβ and Hγ only narrower emission cores are present, and there is no sign of the broad absorption wings of an early B-type main sequence star.

One explanation might be that they are broadened into invisibility by electron scattering. Hamann & Persson (1989) suggested that the very broad wings on the Hα emission line could be explained in that way. They made use of a simple plane-parallel model due to Castor, Smith, & van Blerkom (1970) in which a thin layer of free electrons of optical depth τelies above the region where the line spectrum is produced. The observed spectrum seen from above is the sum, weighted by (1 − e−τe ) and e−τe respectively. Hamann & Persson found that the best match to their Hα profile was obtained with an electron temperature of about 5500 K and τe = 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, namely 5000 K and 0.15, respectively.

With these parameters, as a test a HIRES spectrogram of the narrow-lined B0 V star BD+46◦ 3474 in the 4620 - 4680 A˚ region and near Hα was processed in the same way. Not surprisingly, in the electron-scattering layer the group of strong absorption lines between 4638 and 4654 A˚ was reduced to a broad, shallow blend of central depth Ac = 0.024. But at τe = 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 non-detection of these absorption lines on our HIRES spectrogram of LkHα-101 (which is very noisy in that region), trials show that a τe value of 5 or greater would be required. That would be quite incompatible with the ﬁt to Hα. Therefore the apparent absence of a B-type absorption spectrum cannot be explained in this way.

3.1. The Hydrogen Lines

The Hα emission line is so strong that it was saturated on the ﬁrst two spectrograms of this series. On the third night (2002 December 16) shorter exposures were obtained. Figure 11 shows the spectrum on that date, with the vertical scale expanded for the section longward of Hα. The wing of the Hα emission line extends to at least +1500 km s−1, as noted by Hamann & Persson – 12 –

(1989) and discussed above. The two wings appear quite symmetric when reﬂected about the line center.

The measured equivalent width (W ) of Hα depends critically on how the continuum level is interpolated in from the edges of this very broad emission line. Different trials produce W values that scatter between 500 and 600 A.˚ Table 3 is a collection of published and our own W s. Their dispersion probably reflects the difficulty of consistently defining the continuum level.

The structure near the peak of both Hα and Hβ is shown expanded in Figure 12, both on a velocity scale. The ordinates are ﬂux diﬀerences above the continuum, in units of the continuum. Note that the vertical scales are not the same. The vertical bar at +3.6 km s−1 marks the mean emission line velocity.

Both lines are double-peaked, but not at the velocities or at the separation of the Fe II lines: the Hα separation is 54 km s−1 as compared to about 19 km s−1 for Fe II. The H-line duplicity could be (a) the result of the overlapping of two separate but distinct components (unlikely because there is nothing else in the emission spectrum at those velocities), or (b) a single line that is divided by a near-central reversal produced by higher and cooler H I. That, however, does not explain why the central minima are at diﬀerent velocities (−11km s−1 for Hα, +2km s−1 for Hβ) or why the relative peak heights are not the same, or why Hγ is a single emission line at 0 km s−1.

The H lines present another interesting anomaly. If their strengths are expressed as equivalent widths, the Balmer decrement is very steep: W = 550 ± 50 A˚ for Hα, 39 A˚ for Hβ, 9 A˚ for Hγ7 as compared to the theoretical optically-thin Case B relative intensities of 3.0, 1.0 and 0.46 (for 4 Te = 10 K). If these Balmer lines in LkHα-101 were also optically-thin, the observed decrement can be explained if the emission-line region is more heavily reddened than the continuum to which the W s are referenced. The diﬀerence in AV between the two regions that would reconcile the two decrements depends on the energy distribution of the continuum (for conversion of W s to ﬂuxes) and on the reddening law. For example, if the reddening law were conventional interstellar (RV = 3.1) and if the continuum is that of a 30,000 K blackbody, then ∆AV would be about 1.9 mags. But at the other extreme, if the continuum is produced by dust at 1500 K, then ∆AV would be about 15 mags. Clearly, an AV obtained from H line ratios will be larger than that which measures the foreground extinction.

These considerations would be irrelevant if the Balmer lines are optically-thick, in which case the Balmer decrement could be steeper.

In the radio, the slope of the continuum of the core component of LkHα-101 indicates that it is optically-thick at wavelengths & 0.3 cm (Brown, Broderick, & Knapp 1976). From this a radial α dependence of ne ∝ r can be inferred, and in this way values of α lying between −2 and −2.6

7Cohen & Kuhi (1979) measured the same lines with a spectrum scanner as 587.5, 34.7 and 4.8 A.˚ Hernández et al. (2004) found 464.1, 37.7 and 9.0 respectively, and 6.7 A˚ for Hδ. – 13 – have been obtained (Olnon 1975; Wright & Barlow 1975; Panagia & Felli 1975). Such a radial density profile could be maintained by the outflow of a spherical ionized wind. However, there is no evidence at Hα, or elsewhere in the optical region, of the P Cyg structure that would be expected from such outflow models. No optical counterpart has been found of the absorption component at −400 km s−1 in Brγ reported by Thompson et al. (1976). Nor is any P Cyg structure present in Brγ on a recent high-quality K-band spectrogram at ∼10km s−1 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 α ∼ −2 envelope of LkHα-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 1972). 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, LkHα-101 and NGC 1579 were re-observed 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 A.˚ The nominal dispersion was 90 A˚ mm−1; the actual resolution as measured from a night sky emission line was ∼9 A.˚ The 240′′-long slit was placed through LkHα-101 and rotated so as to cross the brightest section of NGC 1579 to the northeast. The upper panel of Fig. 13 shows the spectrum of LkHα-101 itself, the lower panel the summed spectra of a strip of the nebulosity between 15′′ and 28′′ from the star at P.A. 21◦, both following sky subtraction. The level of the continuous spectrum has been set to 1.0 throughout both spectra. The strong Hβ line is oﬀ scale on both plots, so it has been truncated. The vertical scale of the lower (nebula) plot has been expanded so that Hδ is of the same height in both. The very strong airglow lines at 5577 and 5893 A˚ were only partially removed by the sky subtraction, so they have been blanked out.

To first approximation the spectra of star and nebula are much the same; whether some minor differences between individual emission lines are real would require a detailed investigation on better material. This demonstrates that the spectrum of LkHα-101 as seen directly, and as emitted in another direction and then scattered off the nebula, do not differ significantly. On these much superior spectra no persuasive evidence is seen for the F-type absorption spectrum claimed by Herbig (1956) or by Allen (1972). But neither does one see the interstellar Ca II absorption lines at H and K that one would expect to be present, judging from the strength of interstellar Na I on the HIRES spectra of LkHα-101. Those early low-resolution photographic claims should be dismissed as due to inadequate allowance for sky, or possibly as having been misled by gaps between clumps of emission lines. The presently-available HIRES spectrograms do not extend to the H and K region, and are too noisy at 4300 A˚ to pronounce upon the G band. – 14 –

5. Other Stars in L1482

There are six other fairly bright stars (V = 11 to 13) within about 15′ of LkHα-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 1962-63 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 λ4471/Mg I λ4481 ratio on the HIRES spectrograms, that ratio being type-sensitive in mid- and later B’s. The results are in Table 4.

Star D (= HBC 391) has an interesting spectrum. On a low resolution HARIS spectrogram of the 4000 - 4500 A˚ region the type appears to be early K, but in the red it was classiﬁed by Herbig & Bell (1988) as K7(Li) with weak Hα emission: W (Hα) = 2.3 A.˚ 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”) showing a late-type absorption spectrum and Hα in emission from which they inferred that D is a TTS. The remarkable nature of the spectrum of D becomes apparent at HIRES resolution, which shows that star D is not a conventional TTS. Li I λ6707 is strong (W = 0.26 A).˚ The absorption lines −1 are broad, corresponding to veq sin i about 80km s , but many are asymmetric. The average line centroid is at a velocity of +9 km s−1, but the line minima lie about 15 km s−1 farther longward.

Hα has a most unusual structure: see the upper section of Fig. 14. The two emission components peaking at about −146 and +138 km s−1 appear to be the wings of a single, broad emission line extending to about 300 km s−1 in both directions from the center. The center of this emission line is missing, either because of the superposition of a broad absorption line which itself has an oﬀ-center emission core at about −10km s−1, or because of two separate absorption lines at −44 and +49km s−1, the “emission core” simply being the gap between. The latter may be the correct explanation because Hβ is only a single asymmetric absorption line centered at about +13km s−1, possibly a blend of the two features seen at Hα; no emission is present at Hβ.

The two emission components of Hα have equivalent widths of about 0.34 and 0.16 A,˚ which explains the near-undetectability of Hα emission on the grism spectrograms, especially since at that resolution the emission and central absorption would run together, with a total W of only 0.29 A.˚ However, rough measurement of the Becker & White plot gives W (Hα) ≈ 6 A.˚ Given the unusual nature of the spectrum and the fact that the object is variable at 3.6 cm (Stine & O’Neal 1998) it would not be surprising if the Hα strength varied. According to Preibisch (2002) it is also an X-ray source.

If it can be assumed that an average extinction of about AV = 3.5 mags and a distance modulus of 9.2 mags (both explained below) apply to star D, then MV = +0.5, not far from that of a K-type giant. The spectral similarity is apparent in Fig. 14, where the lower section shows the spectrum – 15 –

−1 of HR 1787, type G9 III-IV, spun up to veq sin i = 80km s . It is plotted as the point D in Fig. 6. Rotational velocities of the B-type stars were estimated by matching the proﬁles of their λ4471 and λ4481 lines with the artiﬁcially spun-up HIRES spectrum of the sharp-lined B8 III star HR 562. The resulting values of veq sin i are in Table 4. The stars’ radial velocities (v⊙) in Table 4 are −1 not determinable with any precision at the higher veq sin i’s; those uncertainties are several km s . The spectrum of star E is not quite normal: lower excitation lines are symmetric but lines of He I, Si II, and Fe III are not, their longward wings being conspicuously steeper than the shortward.

It will be appreciated that in each case these results rest on a single HIRES spectrogram, so that some of the line asymmetries in stars D and E might be due to an imperfectly resolved second spectrum.

The velocities of the ﬁve B-type stars and of star D are near enough to the CO velocity of the L1482 cloud (+6 km s−1) that it is unlikely that they are interlopers from the ﬁeld. Hence they are probably products of the same star-forming activity that has produced LkHα-101 and the cloud’s TTS population. It is interesting that none of these B stars, despite their presumed youth, show any of the peculiarities that would identify them as HAeBe stars.8

6. The Molecular Cloud, and the Distance of NGC 1579

Herbig (1956) originally estimated the distance as about 800 pc on the basis of UBV data and spectral types of two B-type stars that, because they illuminate small reﬂection nebulae, were assumed to be involved in the same cloud as LkHα-101. The types of stars H (the “anon” of Herbig 1971) and E were estimated to be in the range B3 to B5, with E somewhat the earlier, on the basis of low-dispersion Lick photographic spectrograms. Both were therefore assumed to be B4 V, for which MV = −1.5 and (B − V )0 = −0.18 from early compilations by Blaauw and by Johnson (in Strand 1963), so a distance of about 800 pc followed if RV = 3.0. If this calculation is repeated with modern values of the same quantities, the distance becomes about 700 pc. (If the type of star E were actually B2 ± 1, following Becker & White (1988), then its distance would be nearly 1.0 kpc.)

If the same calculation is repeated for the ﬁve B-type stars with their spectral types from the previous section and values for main sequence stars from Schmidt-Kaler (1982), the distances that follow range from 555 pc (star H) to 800 pc (F), with a straight mean of 691 ± 55 pc. Therefore a distance of about 700 pc is reasonable.

8Photometric spectral types can be obtained from the reddening-independent quantity Q = (U-B) - 0.72 (B-V), which is a function of type in the B’s. If the Q’s of Schmidt-Kaler’s luminosity class V are adopted, the resulting photometric types agree with the spectroscopic within one-tenth of a spectral class except for star H, where the photometric type is B6 vs. the spectroscopic B8. The photometric types of stars C and G are B8: and B9, respectively. – 16 –

The line of sight to NGC 1579 passes north of the Tau-Aur clouds, at about 140 pc, and through the Per OB2 association, at roughly 300 pc. Obviously, these distances conﬂict with the 700 pc just inferred. Despite this overlap in the line of sight, CO velocities (Ungerechts & Thaddeus 1987) show that three distinct entities are present. The LSR velocity of the Tau-Aur CO in that general direction is about +6 km s−1, while that of the dense Per OB2 cloud that contains IC 348 and NGC 1333 ranges from +6 to +10 km s−1. NGC 1579 lies on the edge of a broad band of CO that extends from about 4h 40m, +35◦, where it is confused with Per OB2 material, for about 10◦ to the northwest along which the velocity changes from about −1 (near NGC 1579) to −7km s−1. The distinction in velocity between this “northwest feature” and the Tau-Aur, Per OB2 CO is strikingly shown in Ungerecht & Thaddeus’ Figure 3. The LSR velocity of the CO (and of other indicators in Table 6) at NGC 1579 is −1km s−1, indicating that NGC 1579 is associated with the “northwest feature”, not with Tau-Aur or Per OB2.

There is no independent estimate of the distance of the “northwest feature”, but some information can be derived from the duplicity of the cores of the interstellar Na I lines in two of the bright stars embedded in L1482. Fits of two overlapping gaussians to the observed proﬁles produced the velocities given in Table 5. The average LSR velocities of the two Na I components, about −2 and +6km s−1, are near the CO velocities expected of both Per OB2 and “northwest feature” material in that direction. This would indicate that L1482 and NGC 1579 are beyond, not in front of Per OB2, thus supporting a greater distance.

7. The Radio Sources

Becker & White (1988) in a 6 cm survey of the region with the VLA detected nine point sources within about 3′ of LkHα-101, including LkHα-101 itself. Subsequently, Stine & O’Neal (1998) at 3.6 cm with the VLA detected a total of 16 sources in the area, seven being common with the Becker & White list. Stine & O’Neal pointed out that, if they assumed the distance of NGC 1579 is 800 pc, then the radio luminosities of these sources are an order of magnitude higher than those of WTTS in the nearby Tau-Aur clouds, at about 140 pc.

The radio observers, in seeking correspondences with optical or infrared sources, have relied on two papers that announced the detection of an infrared cluster centered on LkHα-101 and gave JHK (and some L) photometry for many stars in the area: Barsony, Schombert, & Kis-Halas (1991) and Aspin & Barsony (1994). We have preferred to use our own catalog (Table 1) of optical and infrared detections. Table 7 summarizes the various correspondences.

Consider ﬁrst the nature of the objects that were detected in radio. Both Becker & White and Stine & O’Neal detected LkHα-101 and star D, but we ﬁnd that optical or infrared point sources are present at only two other radio positions, as follows.

BW6 = SO16 coincides with Table 1/48,49, a close (1′′. 0) double, the brighter (SE) component of which has weak (W (Hα) ≈ 1.7 A)˚ emission. It does not appear in Table 2 because we regard – 17 – grism detections below the 3 A˚ level as unreliable. In this case, the emission is certainly real because the star was already published as HBC 390, with W (Hα) = 2.0 A˚ and a spectral classiﬁcation of M0(Li) from a slit spectrogram reported by Herbig & Bell (1988). Becker & White reproduce a spectrogram (their “lick6”) showing a late-type absorption spectrum with Hα clearly in emission.

BW9 = SO3 coincides with Table 1/66, a star of V = 19.6 without any grism-detectable Hα emission.

But to turn the issue around, note how unsuccessful the radio observations were in detecting known TTS. Table 2 lists 35 Hα emitters (aside from LkHα-101) having W (Hα) ≥ 3 A.˚ Of the 35, 13 have W (Hα) < 10 A˚ and thus would be considered WTTS, which are believed to be more likely than CTTS to be radio emitters (Chiang, Phillips, & Lonsdale 1996). None of those 13 were detected with the VLA.

Therefore, we conclude that the clustering of radio sources around LkHα-101 may very well be real, but if so, they must be heavily obscured pre-main sequence objects of unknown nature. It would be inadvisable to infer a distance to NGC 1579 by comparing the radio ﬂuxes of these sources with those of ordinary WTTS in Tau-Aur.

8. Interstellar Features

The strength of the interstellar Na I lines and diﬀuse interstellar bands (DIBs) in all the NGC 1579 stars observed with HIRES is striking. In fact, some of the stronger DIBs such as λ4428 and λ5780 are obvious on even the low-resolution HARIS spectra. λ6613 appears in Figure 11 of the Hα region of LkHα-101 and in Figure 14 of star D. (In the latter case, it is unusual to see DIBs so clearly against the complex background of a late-type star.) Table 6 gives the W s and velocities for the Na I lines (D1 at 5889 A,˚ D2 at 5895 A)˚ and for all measureable DIBs, and the W s of a sample of the stronger DIBs.

Table 5 gives AV ’s inferred from the (B − V ) colors of the B-type stars, and W s for the same DIBs as observed in the reddened B7 Ia supergiant HD 183143. In the diﬀuse interstellar medium, DIB strengths increase in rough proportion to color excess. Given the usual dispersion in that relationship, the DIBs in these ﬁve B stars are approximately as strong as one might expect, judging from the ratio of their AV ’s to that of HD 183143. If an average foreground AV of about 3.5 mags is applicable to all the brighter stars in L1482, then the additional V extinction of LkHα-101, about 6 - 7 mags, must be local. If essentially all the interstellar features in LkHα-101 are produced in the foreground of L1482, then there is no evidence of any additional interstellar contribution by material in the vicinity of LkHα-101, despite its large local extinction. The explanation may be that the material very near LkHα-101 is depleted in DIB carriers, as has been observed for TTS (Meyer & Ulrich 1984). – 18 –

9. Fine Structure of LkHα-101

As already noted, Danen, Gwinn, & Bloemhof (1995) were able to set an upper limit of 0′′. 34 for the FWHM of the core component of LkHα-101 at 10 µm. Subsequently, Tuthill et al. (2002) resolved the star in the H and K bands at a resolution of about 0′′. 02, using the Keck interferometer. Their images show it as a horseshoe-shaped partial ring about 0′′. 05 (35 AU) across, open on the northeast toward a fainter, bluer companion at a distance of 0′′. 18. They regarded this as providing “deﬁnitive conﬁrmation of the scenario of an accretion disk with a central optically-thin cavity”, the disk being tipped away from the line of sight by . 35◦.

Another interpretation of this structure could be that it represents a tipped annulus either in rotation or expansion around the central star. Assume that the Fe II emission lines in the optical region originate in such a ﬂat annulus of azimuthally uniform surface brightness, and that their duplicity is produced by its rotation. If in Keplerian rotation about a mass M, the circular velocity at r = 0′′. 025 = 17.5 AU (at d = 700 pc) is

1/2 1/2 M 700 pc −1 vc = 27.6 kms . (1) 15 M⊙ d

If the horseshoe structure is really such an annulus with its normal inclined i to the line of sight, then vc will project as a Doppler velocity vp = vc sin i, and the velocity profile will have the form 2 2 −1/2 (vp − v ) . Fig. 15 shows the result of a sample calculation: the annulus was subdivided into 10 sub-annuli between r = 12 and 22 AU, the inclination was 30◦, and the individual line profiles were summed and convolved with the instrumental gaussian. Given the assumptions involved, Fig. 15 is a fair representation of the observed Fe II and Ni II emission line profiles (Fig. 8), but not of those of [O I] and Si II.

On the other hand, the horseshoe might simply be a permanent zone through which ejected material flows, although then one might expect P Cyg structure on the emission lines, which is not observed. If instead it were an expanding ring (which would produce line doubling very much like Fig. 15), at a projected expansion velocity of 10 km s−1, its radius would double in . 9 years. This possibility could be checked by repeating the Tuthill et al. observation some years in the future, but if the structure does change so significantly on so short a time scale it ought to have had some effect on the integrated brightness of LkHα-101. We are not aware of any evidence of long-term variability9.

It is concluded that if the metallic emission lines are indeed produced in the structure reported by Tuthill et al., the rotating annulus explanation appears somewhat the more satisfactory of the two.

9LkHα-101 is listed in the original Catalog of Suspected Variables (as NSV 1618), apparently on the basis of infrared photometry by Strecker & Ney (1974). But those authors state speciﬁcally that the star was not variable over the period of their observations, about 9 months in 1973. – 19 –

10. The Line-of-Sight Structure Near LkHα-101

Fig. 16 contains two images of the immediate vicinity of LkHα-101: left in the K band (∼ 2.2 µm) and right in R (∼ 0.7 µm). Note that neither the “bar” just north of the star or the more distant curved “arc” are seen at R. Fig. 17 is a sketch of an arrangement that can explain this striking difference in the appearance of the nebular structure. It is proposed that, as seen from the earth, LkHα-101 lies behind a 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 the earth that surface is visible only outside the shadow of the foreground cloud. That fraction of the light of LkHα-101 not scattered off the slab penetrates it and illuminates the “bar” and “arc” structures. But to be detected at earth, that radiation must re-emerge 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 Fig. 16.

Becker & White (1988) reproduced a “radiograph” of the immediate region of LkHα-101 constructed from their 6 cm VLA observations (their Figure 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 10′′. It is not the “bar” north of the star seen in Fig. 16, which is only slightly curved, and is about 17′′ 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 LkHα-101.

11. Summary

LkHα-101, which illuminates the reﬂection nebula NGC 1579, has an unusual emission-line spectrum. There is certainly a high-luminosity star in the core of the radio source at that position, whose mass is estimated to be about 15 M⊙ from its Lyc ﬂux (inferred from the radio continuum spectrum), from its position in a V , (V − I) color-magnitude diagram, and the fact that such a mass is compatible with a dynamical interpretation of the duplicity of the star’s Fe II emission lines. The faint Hα-emission stars found in the surrounding molecular cloud (L1482) suggest an age of about 0.5 Myr for that population. Theory suggests that stars of mass &10 M⊙ will complete their pre-main sequence evolution while still heavily obscured, because H-burning in the interior will begin, and the star will arrive on the ZAMS, while heavy accretion is still underway. Therefore not until the foreground material is cleared away will such a star become optically detectable, by which time any T Tauri-like activity should have subsided. We speculate that perhaps there is a time in the early evolution of such a massive ZAMS star when some signature of that recent activity survives, and that in the case of LkHα-101 – which falls in that mass range – its unusual spectrum may be such a signature.

Cohen (1980) envisaged LkHα-101 as a hot star that has ionized the central volume of a thick, dusty circumstellar shell fed by on-going mass outﬂow, from which ultraviolet photons escape – 20 –

(possibly because of a ﬂattened geometry) to ionize the outer envelope. This picture was elaborated by others, notably Simon & Cassar (1984) and Hamann & Persson (1989). Our spectroscopy extends those results and provides new details. Clearly, the emission spectrum of LkHα-101, which must originate near the star, is not produced in a single homogeneous region. That is demonstrated by the variety of 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 §3). This is probably the order in which the density and the level of ionization increase with decreasing distance from the central star. Some of these systematics have been observed in other peculiar stars, so they could indicate an atmospheric structure common to high-luminosity objects of this kind.

However, LkHα-101 does not conform to expectation in several respects. First, if a very young early B-type star is present, no sign of its absorption spectrum can be found in the optical region: the continuum is quite smooth and featureless except for many strong interstellar features. Second, if the material surrounding the core of LkHα-101 is supplied by continuous outflow from a central star, mass-loss rates can be inferred from the radio spectrum via the theory of Wright & Barlow (1975). That theory gives M/v˙ ∞, where v∞ is the terminal velocity of the outflow. Cohen (1980) assumed that v∞ is given by the extended wings of the Hα emission line, and so obtained an M˙ −5 −1 of about 3 × 10 M⊙ yr (alternatively, the Hα wings can be explained by scattering of a core profile in a thin blanket of free electrons: §3.1). Direct evidence of such mass outflow would be P Cyg structure at Hα, but there is no sign of that on the HIRES spectrograms. An absorption component due to cool H I is seen in Hα and Hβ (Fig. 12) but it is near the systemic velocity. The −2 density profile ne ∝ r inferred from the radio spectrum must have another explanation. Third, the high level of ionization in the core, if the radio continuum is interpreted as free- free emission, implies that it is an H II region. But the [O III] lines at 4959, 5007 A˚ ordinarily characteristic of an H II region are not present. One notes that both radio and optical evidence 6 −3 (§3) indicate ne & 10 cm , but the critical density for the upper state of those [O III] transitions is 7 × 105 cm−3 (Osterbrock 1989), so collisional deexcitation would be significant, tending to suppress those lines. But the critical density for the same transitions in [N II] is still lower, so one expects that λλ6548 and 6583 would also not be detectable. Yet there is a line at the position of the stronger of the two, λ6583 (see Fig. 11). Furthermore, the [O II] multiplet at 7318 - 7330 A˚ is present according to Hamann & Persson (1989). Therefore some better explanation for the absence of [O III] is required.

The spectrum of the brightest area of NGC 1579 appears identical to that of LkHα-101, showing that there is no major difference in the spectrum of the source as seen directly or as radiated in another direction and scattered off nearby dust. It would be worthwhile to observe the nebular spectrum at much higher resolution, to see if the finer details are identical.

After LkHα-101, the most luminous stars in the L1482 cloud are ﬁve of types B4 - B9, all of them physically associated with the cloud and members of the cloud cluster on the basis of their agreement in radial velocity with cloud CO. Notably, none appear to be HAeBe stars. If it – 21 –

is assumed that all ﬁve have ZAMS MV ’s and (B − V ) colors, a distance of about 700 pc follows. Thus, although the line of sight passes near the Tau-Aur clouds (at about 140 pc) and through the Per OB2 association (about 300 pc), NGC 1579/L1482 lies well beyond them both. The AV ’s of the B stars average 3.5 mags, as do those of the emission-Hα stars in the cloud for which spectral types are available, although in the latter case there is large dispersion.

If that average AV of 3.5 mags from front and foreground of L1482 also applies to LkHα-101, its additional AV of 6 - 7 mags must originate in the immediate vicinity of the star. Interestingly, since LkHα-101 and all the B-type stars in L1482 show the diﬀuse interstellar band spectrum in comparable strength, and roughly as expected for AV ≈ 3.5 mags, that additional 6 - 7 mags of extinction in front of and around LkHα-101 can contribute little more. It is likely that the circumstellar material, like the dust at many T Tauri stars, is DIB-deﬁcient.

One unusual object in L1482 is the bright star D about 2′ northeast of LkHα-101. It is certainly a member of the cloud cluster because it illuminates its own small patch of reﬂection nebulosity at the edge of NGC 1579, and its radial velocity agrees with those of the B stars. The spectrum is −1 that of a K-type giant, MV about +0.5, broad absorption lines (veq sin i = 80km s ), strong Li I λ6707, and a very complex Hα structure. It is also a radio and an X-ray source. Well above the cluster main sequence, it may be on a radiative track toward the ZAMS at type A.

We are indebted to the National Science Foundation for partial support of this investigation under grants AST 97-30934 and AST 02-04021. Dahm acknowledges support during part of this time from the NASA Graduate Student Researchers Program. We also appreciate unpublished information provided us by Thomas Preibisch and by Ted Simon, and helpful advice from Bo Reipurth. And thanks to Karen Teramura for constructing Figure 17. – 22 –

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

α δ Star (4 +) (35 +) V B − V V − R V − I R R − IJJ − Hb 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 ··· 6.. 29 55.82 16 40.0 ·················· 16.80 0.97 0.76 ··· 7.. 29 55.88 17 37.7 ····················· (18.17) 0.22 ··· 8.. 29 55.96 16 40.6 ·················· 17.45 1.23 0.73 ··· 9.. 29 55.97 17 11.7 ·················· 14.93 1.01 0.58 K7 10. 29 56.35 17 43.0 ·················· 17.56 1.64 1.61 M3 11. 29 56.47 17 51.8 ·················· 14.94 0.93 0.59 ··· 12. 29 56.51 14 29.0 ····················· (16.85) 1.35 M0 13. 29 56.59 16 31.9 ····················· (17.19) 0.09 ··· 14. 29 57.03 19 32.7 18.65 2.38 1.52 3.13 17.13 1.61 ············ 15. 29 57.22 19 18.2 ············ 20.96 2.63 ············ 16. 29 57.24 15 02.7 ····················· (16.41) 1.90 ··· 17. 29 57.32 19 29.9 22.75 ··· 1.31 2.74 21.44 1.43 ············ 18. 29 57.44 19 24.1 16.45 2.05 1.36 2.73 15.09 1.38 ············ 19. 29 57.58 17 37.6 ·················· 16.67 1.28 0.83 ··· 20. 29 57.62 14 33.8 ·················· 18.46 2.45 1.01 ··· 21. 29 58.24 19 48.7 22.41 ··· 1.27 2.64 21.13 1.37 ············ 22. 29 58.25 15 35.3 ·················· 17.47 1.69 1.10 ··· 23. 29 58.36 16 44.0 ·················· 17.43 1.21 0.39 ··· 24. 29 58.40 19 26.1 22.54 ··· 1.96 4.62 20.59 2.66 ············ 25. 29 58.51 16 52.9 21.74 ··· 1.58 3.19 20.16 1.61 15.95 1.01 0.45 ··· 26. 29 58.60 16 30.2 22.08 ··· 1.28 2.93 20.79 1.64 16.12 1.16 0.49 ··· 27. 29 58.61 16 17.4 ············ 20.99 3.03 ········· M2 28. 29 58.67 16 41.8 ·················· 16.26 1.13 0.65 ··· 29. 29 58.68 16 17.8 ·················· 14.37 1.26 0.78 ··· 30. 29 59.19 18 48.6 ············ 21.80 3.00 ········· K9 31. 29 59.38 13 07.5 17.01 1.88 1.47 3.21 15.55 1.74 ············ 32. 29 59.72 13 34.3 17.35 2.51 1.62 3.40 15.73 1.79 ············ 33. 29 59.89 18 50.7 20.24 2.65 1.88 3.72 18.35 1.84 ············ 34. 29 59.94 15 15.3 ·················· 17.19 1.48 1.09 ··· 35. 29 59.99 18 20.3 22.09 ··· 1.94 4.36 20.15 2.42 ············ 36. 30 00.02 14 17.4 ······························ 37. 30 00.04 16 55.1 21.43 2.05 1.34 2.75 20.09 1.41 16.40 1.06 0.54 ··· 38. 30 00.15 17 15.2 21.48 2.38 1.82 4.16 19.67 2.34 14.61 0.96 0.49 M2 39. 30 00.33 19 04.4 22.29 ··· 1.78 3.80 20.51 2.01 ············ 40. 30 00.40 12 39.5 14.63 1.06 0.63 1.22 14.00 0.59 ············ 41. 30 00.47 16 35.3 ·················· 18.11 1.41 0.44 ··· 42. 30 00.51 18 17.0 19.05 1.94 1.30 2.73 17.76 1.43 ············ 43. 30 00.53 19 25.9 21.25 2.26 1.32 3.25 19.93 1.93 ············ 44. 30 00.63 17 18.4 19.76 2.50 1.68 4.10 18.08 2.43 12.89 1.01 0.48 M2(Li) 45. 30 00.66 16 40.4 ························ 0.80 ··· – 27 –

Table 1—Continued

α δ Star (4 +) (35 +) V B − V V − R V − I R R − IJJ − Hb H − K SpT

46. 30 00.76 17 57.7 ·················· 17.80 1.85 1.62 ··· 47. 30 01.08 15 32.3 ·················· 17.27 1.88 1.09 ··· 48. 30 01.16 17 26.0 18.19 2.38 1.65 ··· 16.54 ··············· 49. 30 01.19 17 24.7 16.28 2.09 1.39 3.15 14.89 1.77 ········· K7(Li) 50. 30 01.24 14 29.2 ·················· 18.82 2.67 1.74 ··· 51. 30 01.30 14 23.0 22.27 ··· 1.85 4.59 20.42 2.74 14.94 1.02 0.44 ··· 52. 30 01.38 15 02.3 ·················· 17.02 2.07 1.20 ··· 53. 30 01.49 16 12.1 20.43 2.08 1.25 2.80 19.17 1.55 15.10 0.92 0.57 ··· 54. 30 01.58 16 04.9 ·················· 16.80 1.28 0.76 ··· 55. 30 01.62 15 40.3 ····················· (17.38) 0.66 ··· 56. 30 01.79 15 32.2 ····················· (17.43) 0.81 ··· 57. 30 01.88 14 12.7 ····················· (17.14) 1.03 ··· 58. 30 01.88 15 51.2 ············ 22.29 3.26 15.27 1.22 0.79 M3 59. 30 01.92 17 07.8 ············ 21.83 1.23 ············ 60. 30 01.97 16 44.5 21.76 ··· 1.39 3.92 20.38 2.53 15.52 0.74 0.50 ··· 61. 30 02.14 19 40.8 22.49 ··· 1.49 3.04 20.99 1.55 ············ 62. 30 02.21 14 28.3 ····················· (16.90) 1.01 ··· 63. 30 02.21 17 16.8 ····················· (17.09) 0.86 ··· 64. 30 02.22 16 34.2 22.39 ··· 1.61 4.26 20.78 2.65 15.41 0.98 0.61 M3 65. 30 02.57 18 08.6 ············ 22.04 3.25 ············ 66. 30 02.64 15 14.9 19.60 2.82 1.95 4.21 17.65 2.26 12.17 1.41 0.69 ··· 67. 30 02.71 17 39.7 ············ 21.67 3.01 15.69 0.92 0.55 ··· 68. 30 02.88 18 44.4 ············ 21.12 1.92 ············ 69. 30 02.89 14 41.2 21.21 ··· 2.05 4.45 19.16 2.40 13.26 1.50 0.63 ··· 70. 30 03.20 14 21.5 ············ 21.69 3.20 14.95 0.99 0.63 M3 71. 30 03.31 15 44.2 19.75 2.38 1.66 3.42 18.09 1.77 13.56 1.18 0.51 ··· 72. 30 03.58 16 38.0 19.93 2.54 1.74 3.77 18.19 2.04 13.54 1.12 0.56 M2(Li) 73. 30 03.66 14 20.9 ····················· (16.34) 3.05 ··· 74. 30 03.78 15 59.1 ····················· (17.10) 1.01 ··· 75. 30 03.80 14 36.2 22.25 ··· 1.35 2.82 20.90 1.47 ··· (16.74) 1.43 ··· 76. 30 03.95 15 10.2 ····················· (16.92) 0.73 ··· 77. 30 03.96 16 25.4 19.35 2.16 1.38 2.83 17.97 1.45 14.18 0.96 0.47 ··· 78. 30 04.16 16 27.5 22.21 ··· 2.11 4.80 20.10 2.68 14.01 1.27 0.76 M3 79. 30 04.20 17 27.5 21.66 ··· 1.71 4.12 19.95 2.41 14.75 1.15 0.50 ··· 80. 30 04.32 15 57.5 ·················· 17.94 1.59 0.68 ··· 81. 30 04.40 12 39.5 22.76 ··· 2.04 4.10 20.72 2.06 ············ 82. 30 04.59 16 04.4 ············ 21.62 2.79 14.74 1.39 0.90 ··· 83. 30 04.62 15 01.6 ············ 22.08 3.49 14.77 0.95 0.70 M2 84. 30 04.62 15 24.2 ·················· 16.49 0.93 0.54 ··· 85. 30 04.80 17 22.4 ·················· 16.34 1.15 0.70 ··· 86. 30 05.28 19 26.3 22.31 ··· 1.56 2.89 20.75 1.33 ············ 87. 30 05.40 15 32.5 ·················· 15.94 1.22 0.67 ··· 88. 30 05.48 17 52.3 21.64 ··· 1.70 3.27 19.94 1.57 15.94 1.28 0.42 K9 89. 30 05.52 17 08.2 22.39 ··· 1.69 3.34 20.70 1.64 15.95 0.67 0.52 ··· 90. 30 05.55 18 07.6 20.30 2.16 1.41 2.89 18.89 1.48 ············ – 28 –

Table 1—Continued

α δ Star (4 +) (35 +) V B − V V − R V − I R R − IJJ − Hb H − K SpT

91. 30 05.62 16 40.6 ·················· 15.55 0.56 ······ 92. 30 05.63 17 06.6 21.54 1.73 1.64 3.17 19.90 1.53 ············ 93. 30 05.78 16 25.6 ·················· 17.61 1.18 0.55 ··· 94. 30 05.83 15 23.1 ·················· 18.12 1.53 0.69 ··· 95. 30 05.89 17 00.7 21.86 ··· 2.16 4.86 19.70 2.70 13.71 0.74 0.69 M3 96. 30 05.94 17 02.7 ····················· (16.01) 0.76 ··· 97. 30 06.12 17 07.3 22.43 ··· 1.61 2.98 20.82 1.37 16.76 1.11 0.83 ··· 98. 30 06.26 18 42.8 22.37 ··· 1.57 3.26 20.80 1.68 ········· K8 99. 30 06.40 16 45.9 22.86 ··· 2.29 5.15 20.58 2.86 14.31 0.75 0.56 ··· 100 30 06.65 17 53.1 21.69 ··· 1.50 3.72 20.19 2.22 14.44 1.18 0.68 M3 101 30 06.66 15 07.7 ·················· 18.22 1.26 0.52 ··· 102 30 07.04 17 41.9 22.44 ··· 1.80 3.32 20.64 1.52 16.68 1.19 ··· K7 103 30 07.25 15 17.7 ····················· (17.68) 1.24 ··· 104 30 07.40 14 35.6 ·················· 17.46 0.49 0.12 ··· 105 30 07.43 14 58.6 18.32 2.31 1.61 3.41 16.71 1.80 12.13 1.56 ··· K8(Li) 106 30 07.44 19 49.2 18.74 2.20 1.53 3.19 17.20 1.66 ············ 107 30 07.50 17 54.4 20.19 2.56 1.87 3.94 18.32 2.07 13.31 1.31 0.62 ··· 108 30 07.52 15 59.5 21.43 ··· 1.81 3.97 19.62 2.16 14.63 1.26 0.52 ··· 109 30 07.71 15 10.9 20.57 2.00 1.75 3.93 18.82 2.18 14.04 1.08 0.41 ··· 110 30 07.73 17 22.0 ·················· 16.45 0.87 0.44 ··· 111 30 07.75 15 49.0 20.24 2.23 1.86 4.00 18.38 2.14 13.41 1.34 0.89 ··· 112d 30 07.82 14 09.7 19.53 2.17 1.69 3.73 17.84 2.04 13.00 0.81 0.42 M2(Li) 113 30 07.93 15 12.0 ····················· (17.80) 0.06 ··· 114 30 08.01 16 50.2 ·················· 14.86 1.34 0.80 ··· 115 30 08.21 17 26.5 ····················· (17.43) 0.81 ··· 116 30 08.23 17 23.8 ·················· 15.20 1.13 0.84 ··· 117 30 08.24 14 10.7 ·················· 17.41 2.52 1.56 ··· 118 30 08.36 14 39.8 19.45 2.46 1.82 4.15 17.63 2.33 12.27 1.52 0.64 M2(Li) 119 30 08.50 14 48.1 ····················· (16.14) 0.69 ··· 120 30 08.56 15 31.5 ············ 21.14 2.89 14.98 0.97 0.58 ··· 121 30 08.71 14 23.3 ·················· 15.53 1.45 0.93 ··· 122 30 08.74 14 38.3 19.00 2.81 1.85 3.96 17.15 2.11 12.05 ········· 123 30 08.78 13 54.0 ············ 21.32 1.96 ············ 124 30 08.82 16 09.5 19.87 2.31 1.82 4.18 18.05 2.35 12.85 1.19 0.49 ··· 125 30 08.87 15 00.0 ·················· 17.91 1.64 0.85 ··· 126 30 08.97 14 33.3 21.63 ··· 2.12 4.90 19.51 2.79 13.13 1.51 0.79 M4 127 30 08.98 15 10.2 ············ 20.70 1.66 16.11 1.27 0.52 ··· 128 30 09.14 17 20.8 21.66 ··· 1.72 4.15 19.94 2.43 14.68 1.22 0.52 ··· 129 30 09.22 17 26.7 22.09 ··· 1.42 3.16 20.67 1.74 15.97 1.38 0.50 ··· 130 30 09.22 17 39.4 ·················· 17.28 1.05 0.91 ··· 131 30 09.38 17 11.7 ·················· 17.31 1.21 0.50 ··· 132 30 09.43 17 41.0 22.34 ··· 2.01 4.87 20.33 2.86 13.75 1.01 0.69 M3 133 30 09.51 14 41.1 ············ 22.17 2.59 16.00 1.84 1.21 ··· 134 30 09.65 16 58.3 ·················· 17.53 1.16 0.65 ··· 135 30 09.71 17 06.0 ·················· 18.17 1.46 0.87 ··· – 29 –

Table 1—Continued

α δ Star (4 +) (35 +) V B − V V − R V − I R R − IJJ − Hb H − K SpT

136 30 09.74 14 33.2 ·················· 18.16 2.62 1.13 ··· 137 30 09.78 15 19.9 ············ 21.81 2.52 16.76 0.78 0.49 ··· 138 30 09.79 13 52.8 21.95 ··· 1.96 4.66 19.99 2.70 ············ 139 30 09.85 14 17.1 21.97 ··· 2.14 4.02 19.83 1.88 15.12 1.48 1.07 K2 140e 30 09.92 15 54.7 18.85 2.43 0.00 4.19 18.85 4.19 12.01 0.65 0.75 ··· 141 30 09.97 15 38.4 ············ 19.97 2.65 13.75 1.37 0.94 ··· 142 30 10.03 14 23.2 20.40 2.89 1.92 4.07 18.48 2.15 13.35 1.33 0.55 ··· 143 30 10.08 14 50.0 20.95 2.08 1.17 2.97 19.79 1.81 16.51 0.65 0.30 ··· 144 30 10.20 17 55.3 ····················· (16.59) 0.82 ··· 145 30 10.28 17 07.8 ·················· 16.21 1.16 0.71 ··· 146 30 10.29 15 56.5 ············ 19.43 2.64 13.32 0.53 0.56 M3(Li) 147 30 10.45 16 25.6 ····················· (17.12) 2.10 ··· 148 30 10.47 19 43.8 21.34 ··· 1.20 2.89 20.14 1.68 ············ 149 30 10.53 19 06.5 20.23 ··· 1.47 3.61 18.76 2.15 ············ 150 30 10.57 16 50.3 ············ 19.27 2.36 12.95 1.67 1.07 ··· 151 30 10.59 16 56.2 ············ 18.81 2.18 13.71 1.14 ······ 152 30 10.72 14 55.6 ·················· 17.84 1.10 0.65 ··· 153 30 10.87 17 11.8 ····················· (17.08) 0.51 ··· 154 30 10.89 16 13.3 ·················· 15.87 1.13 0.81 ··· 155 30 10.94 16 21.3 ·················· 16.35 1.01 1.13 ··· 156 30 10.96 13 40.9 ············ 21.53 3.26 ············ 157 30 11.08 16 04.0 ·················· 15.65 2.18 1.95 K6 158 30 11.20 19 41.0 ············ 21.44 2.18 ············ 159 30 11.26 14 49.6 ····················· (17.01) 1.45 ··· 160 30 11.30 16 13.4 ·················· 14.11 1.54 0.79 ··· 161 30 11.52 17 02.7 18.57 ··· 1.64 3.45 16.93 1.80 12.21 1.16 ······ 162 30 11.53 14 13.4 13.28 0.89 ············ 11.45 0.46 0.06 F9 163 30 11.59 14 30.5 ·················· 17.78 1.62 0.82 ··· 164 30 11.64 16 52.8 ·················· 16.65 0.97 1.07 M0 165 30 11.67 17 04.0 ·················· 15.61 ········· 166 30 11.74 16 53.9 ····················· (17.29) 0.81 ··· 167 30 11.76 16 31.7 ·················· 13.88 1.74 1.14 ··· 168 30 12.19 14 51.0 ·················· 16.48 1.64 1.07 ··· 169 30 12.22 19 53.6 ············ 21.28 2.94 ············ 170 30 12.23 15 47.3 ·················· 15.02 1.36 1.01 ··· 171 30 12.24 18 40.2 15.13 1.58 0.94 ··· 14.19 ··············· 172 30 12.32 19 34.1 20.12 1.54 1.02 2.16 19.09 1.13 ············ 173 30 12.43 16 28.4 ·················· 13.59 1.70 1.19 ··· 174 30 12.55 16 00.9 19.17 ··· 1.00 2.69 18.17 1.69 14.68 0.70 0.37 ··· 175 30 12.77 17 21.3 ·················· 14.02 1.32 0.98 ··· 176 30 12.88 15 17.4 21.55 ··· 1.87 4.32 19.69 2.46 14.38 0.22 0.41 M3 177 30 12.98 15 11.6 ····················· (17.37) 0.98 ··· 178 30 13.01 16 33.3 ·················· 14.79 1.06 1.07 ··· 179 30 13.01 17 25.4 ·················· 15.67 1.44 0.68 ··· 180 30 13.05 13 59.5 13.85 1.39 0.88 ··· 12.97 ············ F6 – 30 –

Table 1—Continued

α δ Star (4 +) (35 +) V B − V V − R V − I R R − IJJ − Hb H − K SpT

181 30 13.08 15 18.8 ·················· 17.69 0.94 1.23 ··· 182f 30 13.09 16 31.2 ·················· 15.24 0.00 0.44 ··· 183 30 13.10 16 02.1 ·················· 17.35 1.63 0.63 ··· 184 30 13.17 16 33.6 ·················· 16.29 0.29 0.68 ··· 185 30 13.30 18 43.4 ············ 21.72 1.22 ············ 186 30 13.40 15 56.3 ·················· 14.98 1.90 1.17 ··· 187 30 13.41 18 11.4 19.84 2.84 1.75 4.21 18.09 2.46 ········· M4(Li) 188 30 13.44 15 41.5 ·················· 17.08 1.21 0.89 ··· 189 30 13.70 15 26.3 ·················· 16.04 0.74 0.51 M3 190 30 13.82 16 21.0 ·················· 15.57 2.28 ······ 191 30 14.07 15 22.7 18.09 1.73 1.10 2.11 16.99 1.02 14.57 0.71 0.26 ··· 192 30 14.26 17 51.9 22.06 ··· 1.86 4.86 20.20 2.99 14.31 0.78 0.59 M4 193 30 14.35 17 48.2 ····················· (16.53) 1.53 ··· 194g 30 14.44 16 24.5 15.71 2.20 2.38 ··· 13.33 ··············· 195 30 14.61 15 59.8 ····················· (17.10) 2.43 ··· 196 30 14.62 16 12.9 ·················· 16.33 ········· 197 30 15.17 15 30.6 ·················· 17.31 1.25 0.82 ··· 198 30 15.20 16 40.4 ·················· 12.28 1.01 1.10 ··· 199 30 15.33 16 40.1 ····················· (14.70) 1.13 ··· 200 30 15.27 16 08.3 ·················· 16.55 2.61 1.66 ··· 201 30 15.44 16 33.3 ·················· 15.84 2.12 1.31 ··· 202 30 15.45 14 23.0 ·················· 18.15 1.52 0.70 ··· 203 30 15.51 15 16.5 ·················· 18.53 1.12 0.36 ··· 204 30 15.61 16 00.3 ·················· 16.31 2.17 1.28 ··· 205 30 15.64 17 38.4 ············ 21.18 2.63 14.05 1.76 1.35 M1 206 30 15.67 13 29.3 ············ 21.80 2.35 ············ 207 30 15.88 15 27.4 ············ 20.45 1.71 15.28 1.23 0.70 ··· 208 30 15.95 14 34.4 ····················· (17.50) 0.98 ··· 209 30 15.97 19 46.4 20.67 2.22 1.50 3.73 19.17 2.23 ············ 210 30 16.05 17 27.5 ············ 20.97 2.56 14.33 1.60 0.88 M2 211 30 16.11 16 10.0 ·················· 15.98 2.11 1.37 ··· 212 30 16.30 15 24.7 ············ 21.10 2.76 14.12 1.35 0.91 ··· 213 30 16.40 19 24.9 20.57 3.39 1.38 3.12 19.19 1.74 ············ 214 30 16.46 14 38.9 ·················· 16.30 1.08 0.82 ··· 215 30 16.56 15 42.7 ············ 19.40 2.27 13.68 1.39 0.70 M3 216 30 16.60 19 23.8 17.41 2.09 1.32 2.82 16.09 1.50 ············ 217 30 16.69 15 55.7 ·················· 15.18 1.29 0.81 ··· 218 30 16.78 15 18.5 ····················· (17.86) 0.06 ··· 219 30 16.79 14 45.2 ············ 22.42 3.60 15.23 1.39 0.86 ··· 220 30 16.80 15 26.2 ····················· (16.44) 1.00 ··· 221 30 17.00 16 47.1 ·················· 15.41 1.82 1.05 ··· 222 30 17.13 16 16.3 ·················· 13.99 2.21 1.40 ··· 223 30 17.14 14 56.1 21.11 ··· 2.19 4.40 18.91 2.21 13.16 1.42 0.65 ··· 224 30 17.23 14 36.7 ·················· 17.59 1.45 0.46 ··· 225 30 17.24 15 38.8 17.67 2.11 1.57 3.27 16.10 1.70 11.81 1.15 0.65 K8(Li) – 31 –

Table 1—Continued

α δ Star (4 +) (35 +) V B − V V − R V − I R R − IJJ − Hb H − K SpT

226 30 17.25 16 03.8 ·················· 16.12 1.73 1.19 ··· 227 30 17.28 19 17.5 19.28 2.96 1.82 3.81 17.46 1.99 ············ 228 30 17.37 15 21.4 ·················· 17.90 1.19 0.78 ··· 229 30 17.55 16 26.6 ·················· 16.65 1.82 1.52 ··· 230 30 17.74 17 13.7 ·················· 17.12 2.15 1.59 ··· 231 30 17.74 18 11.4 ············ 21.68 2.15 ············ 232 30 17.92 16 08.4 ·················· 17.24 2.42 1.58 ··· 233 30 18.06 18 18.8 22.60 ··· 1.85 4.45 20.75 2.60 ········· M4 234 30 18.43 17 09.1 ····················· (16.97) 1.25 ··· 235 30 18.53 15 50.4 ····················· (16.97) 1.21 ··· 236 30 18.68 16 42.9 ·················· 15.24 1.56 1.22 ··· 237 30 18.80 16 41.9 ·················· 15.30 1.78 1.64 ··· 238 30 18.82 15 54.7 ·················· 17.29 1.94 1.20 ··· 239 30 18.89 17 12.1 ·················· 18.13 1.26 0.82 ··· 240 30 19.01 17 36.2 ·················· 16.90 1.75 0.94 ··· 241h 30 19.15 17 45.9 13.43 1.87 ····················· K0: 242 30 19.24 16 02.5 ·················· 16.09 2.66 1.54 ··· 243 30 19.35 14 00.7 ············ 20.81 2.74 ········· M3 244 30 19.39 15 57.3 ·················· 15.39 2.70 1.77 ··· 245 30 19.45 16 58.0 ····················· (16.93) 1.02 ··· 246 30 19.48 17 24.4 21.14 ··· 1.76 4.61 19.37 2.85 13.50 1.10 0.56 ··· 247 30 19.46 16 34.9 ·················· 16.54 2.90 1.84 ··· 248 30 19.49 17 42.0 ····················· (15.36) 0.67 ··· 249 30 19.50 15 50.9 21.42 ··· 1.63 3.93 19.79 2.30 15.10 0.86 0.39 ··· 250 30 19.54 17 03.1 ····················· (17.33) 1.01 ··· 251 30 19.64 17 00.2 ·················· 15.90 1.85 0.99 ··· 252 30 19.65 18 17.1 20.38 2.38 1.70 4.50 18.69 2.80 ············ 253 30 19.79 14 21.9 21.85 ··· 1.78 4.26 20.07 2.48 14.69 1.21 0.69 M3 254 30 19.89 15 11.2 21.60 ··· 1.81 4.27 19.79 2.47 14.23 1.18 0.51 ··· 255 30 19.94 15 40.7 ····················· (18.13) 0.59 ··· 256 30 19.94 18 14.9 ············ 21.94 1.85 ············ 257 30 19.95 14 38.7 ·················· 15.36 1.57 0.82 ··· 258 30 20.18 17 35.7 ·················· 17.91 1.05 0.45 ··· 259 30 20.21 16 55.1 ·················· 16.56 1.98 1.10 ··· 260 30 20.41 14 56.0 ·················· 15.45 1.57 0.67 M0 261 30 20.77 15 12.8 ·················· 16.39 1.75 0.90 ··· 262 30 20.86 15 05.4 22.29 ··· 1.60 3.82 20.69 2.23 16.43 0.79 0.29 M3 263 30 20.93 13 03.9 20.60 2.56 1.56 4.08 19.03 2.51 ············ 264 30 20.94 15 24.0 ·················· 16.23 1.69 0.78 ··· 265 30 21.00 15 40.3 ············ 22.41 2.45 16.72 1.31 0.55 ··· 266 30 21.63 18 15.2 22.10 ··· 1.47 4.10 20.63 2.64 ············ 267 30 21.74 14 54.2 ·················· 15.63 1.55 0.65 ··· 268 30 22.19 14 53.1 ····················· (17.31) 0.99 ··· 269 30 22.33 16 47.6 ·················· 17.72 1.17 0.61 ··· 270 30 22.34 16 32.8 ·················· 16.24 1.39 0.69 ··· – 32 –

Table 1—Continued

α δ Star (4 +) (35 +) V B − V V − R V − I R R − IJJ − Hb H − K SpT

271 30 22.83 17 00.0 21.84 ··· 1.88 3.85 19.96 1.97 14.96 1.33 0.50 ··· 272i 30 22.97 18 16.2 ··· (17.86) ······ 14.46 ············ B8 273 30 22.98 15 10.8 21.92 ··· 1.31 3.39 20.61 2.08 16.43 0.82 0.29 ··· 274 30 23.02 15 03.3 ·················· 18.59 1.67 0.73 ··· 275 30 24.06 18 12.2 21.14 1.82 1.54 2.68 19.59 1.13 ············ 276 30 24.35 17 30.2 22.70 ··· 1.66 3.50 21.03 1.84 16.57 0.76 0.50 ··· 277 30 25.06 18 17.4 21.86 1.78 1.27 2.48 20.59 1.22 ············ 278 30 25.10 13 51.4 18.77 2.48 1.68 3.53 17.09 1.85 ············ 279 30 25.48 18 12.5 ············ 21.69 1.14 ············ 280 30 25.64 14 42.4 ·················· 16.84 1.68 0.97 ··· 281 30 25.83 18 23.7 21.95 1.53 1.22 2.32 20.74 1.10 ············ 282 30 26.57 19 41.1 ············ 21.03 1.64 ············ 283 30 26.93 18 07.2 19.31 2.57 1.49 3.09 17.82 1.60 ············ 284 30 26.96 14 49.0 ·················· 17.63 1.20 0.90 ··· 285 30 27.08 16 18.7 ·················· 17.97 1.05 0.51 ··· 286 30 27.56 14 15.5 ·················· 17.71 1.91 1.00 ··· 287 30 27.85 18 40.2 ············ 21.81 1.53 ············ 288 30 27.99 15 15.7 21.88 ··· 1.61 4.34 20.27 2.73 14.97 0.65 0.52 ··· 289 30 27.99 16 43.3 ············ 20.81 1.65 16.42 0.98 0.38 ··· 290 30 28.04 17 11.5 ····················· (17.95) 0.55 ··· 291 30 28.08 18 32.7 22.19 ··· 0.98 2.13 21.21 1.15 ············ 292 30 28.54 15 51.1 22.62 ··· 1.53 4.00 21.09 2.47 16.64 0.47 0.46 ··· 293 30 28.58 15 35.5 21.94 ··· 2.19 4.32 19.74 2.13 14.32 1.25 0.59 ··· 294 30 28.73 17 25.6 20.72 2.73 1.81 3.53 18.91 1.73 14.39 0.99 0.53 ··· 295 30 28.82 16 02.4 ············ 21.37 2.20 16.11 1.21 0.55 ··· 296 30 29.29 16 33.1 ·················· 17.49 0.68 ······ 297 30 29.49 17 16.3 ·················· 18.78 1.60 0.83 ··· 298 30 29.64 16 35.5 23.36 ··· 2.15 3.79 21.21 1.64 17.23 1.06 0.22 ··· 299 30 29.74 19 05.9 21.52 ··· 1.71 3.48 19.81 1.77 ············ 300 30 30.16 14 27.9 ·················· 15.55 0.84 0.49 ··· 301 30 30.27 15 21.3 ····················· (17.59) 0.57 ··· 302 30 30.32 14 29.9 ·················· 17.33 1.53 0.82 ··· 303 30 30.41 18 34.4 22.05 ··· 1.86 4.51 20.19 2.66 ············ 304 30 30.49 17 45.5 21.74 ··· 1.66 3.72 20.07 2.06 13.96 1.38 0.78 ··· 305 30 30.62 15 34.2 ····················· (17.09) 0.98 ··· 306 30 30.69 17 43.3 21.48 ··· 1.80 4.53 19.68 2.73 14.08 1.11 0.46 ··· 307 30 30.71 17 38.7 ············ 21.71 1.68 ··· (16.97) 0.20 ··· 308 30 31.03 15 31.8 ············ 21.09 2.07 15.79 1.24 0.54 ··· 309 30 31.27 15 15.4 ····················· (17.93) 1.31 ··· 310 30 31.30 14 19.1 ····················· (16.85) 0.75 ··· 311 30 31.52 17 56.7 19.63 2.28 1.63 3.29 18.00 1.66 13.77 1.18 0.37 ··· 312 30 31.70 14 23.9 13.51 1.09 0.64 1.23 12.87 0.58 11.29 ········· 313 30 31.72 18 42.6 19.18 1.85 1.18 2.39 18.00 1.20 ············ 314 30 31.77 15 36.3 ·················· 17.80 1.19 0.36 ··· 315 30 31.94 15 37.1 ·················· 18.58 1.02 0.41 ··· – 33 –

Table 1—Continued

α δ Star (4 +) (35 +) V B − V V − R V − I R R − IJJ − Hb H − K SpT

316 30 31.99 16 45.0 18.76 1.97 1.11 3.03 17.65 1.91 14.03 0.79 0.23 ··· 317 30 32.04 14 37.0 ·················· 16.37 1.78 0.73 ··· 318 30 32.05 15 39.0 ····················· (18.27) 0.50 ··· 319 30 32.37 16 15.7 21.46 1.86 1.38 2.80 20.09 1.43 16.41 0.50 0.12 ··· 320 30 32.42 18 27.0 22.22 ··· 1.58 2.92 20.64 1.34 ············ 321 30 33.16 17 03.7 ·················· 16.75 0.78 0.27 ··· 322 30 33.30 16 46.7 ·················· 17.33 1.13 0.30 ··· 323 30 33.83 16 50.2 ·················· 16.43 1.20 0.54 ··· 324 30 34.38 15 20.3 ·················· 17.71 ········· 325 30 34.44 15 02.8 ·················· 17.20 ·········

aRight ascension (α) is in units of hours, minutes, and seconds. Declination (δ) is in units of degrees, arcminutes, and arcseconds. The BV RI colors have been converted to the Johnson-Kron-Cousins photometric system, whereas the JHK colors are on the new Mauna Kea ﬁlter system. bWhen only H and K measurements are available, this column contains the H magnitude in parentheses. cIntegrated magnitudes of an unresolved binary. dEmission lines including [O I] λλ6300, 6363, [N II] λ6584, and [S II] λλ6716, 6730 are present. eIntegrated magnitudes of an unresolved triple. f The IR colors of this star are severely aﬀected by the steep brightness gradient of the LkHα-101 infrared nebulosity; its J and H magnitudes are particularly uncertain. gThis is LkHα-101. Two stars are present within the glare of this star: one to the northeast at α = 4 30 14.76, δ = +35 16 34.5 (J=14.9, (J − H)=1.7, (H − K)=1.0) and one to the southwest at α = 4 30 13.73, δ = +35 16 20.9 (J=16.1, (J − H)=2.2, (H − K)=1.6). Because of the glare from LkHα-101, the photometry for these objects is rather uncertain. hThis is star D. The spectral type is uncertain, but likely K0 III. See the text for details. iThis is a double star, and is saturated in V . Therefore, the B magnitude is in parentheses in the (B − V ) column. – 34 –

Table 2. Hα Emission Stars in NGC 1579

IHα Table 1 W (Hα)a

187.... 10..... 8 188.... 27..... 22 189.... 30..... 79 190.... 32..... 3 191.... 44..... 4 192.... 63..... 4 193.... 70..... 9 194.... 72..... 8 195.... 78..... 54, 16 196.... 83..... 26 197.... 95..... 94, 32 198.... 100.... 61 199.... 105.... 20, 11 200.... 107.... 6 201.... 111.... 69 202.... 112.... 75, 35 203.... 118.... 4 204.... 122.... 10 205.... 126.... 34 206.... 132.... 21 207.... 139.... 224 208.... 140.... 33 209.... 151.... 51 210.... 157.... 372 211.... 180.... 6 212.... 187.... 6, 7 213.... 192.... 8 214.... 194.... 1075 215.... 205.... 45 216.... 215.... 24 217.... 225.... 3, 3 218.... 233.... 6 219.... 243.... 182, 127 220.... 253.... 176, 102 221.... 303.... 164 222.... 304.... 20

aAll equivalent widths are in A.˚ Those with two measurements have the grism value ﬁrst and the GMOS value second. – 35 –

Table 3. LkHα-101 Hα Equivalent Widths

Date W (A)˚ Reference

1970-71 >400 Herbig (1971) 1976 588 Cohen & Kuhi (1979) 1983 1050 Hamann & Persson (1989) 1999-2000 464 Hern´andez et al. (2004) 2002 Dec. 550 ± 50 HIRES, this paper 2003 Jan. 1075 WFGS, this paper – 36 –

Table 4. Coordinates and Lick UBV Photometry of Bright Stars in L1482a

Star α δ V (B − V ) (U − 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 ∗ Eb 4 30 10.9 +35 19 23 10.96 +0.87 +0.10 B4 +12 90 ∗ Dc 4 30 19.2 +35 17 46 13.23 +1.84 +1.08: K0: +9 80 ∗ Hd 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 ········· ∗ 101f 4 30 14.4 +35 16 24 15.71 +2.20 ············

aThese UBV observations were obtained at Lick Observatory by B. Paczynski on 1962 De- cember 26 and 1963 January 2. Those stars that illuminate individual reﬂection nebulae are marked with asterisks. The coordinates are in standard units, equinox J2000. The velocities quoted in the last two columns are given in units of km s−1. A colon after a number marks a particularly uncertain value. bHD 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. cThis 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. dThe nebulosity at H is IC 2067. eOur values (Table 1/40) are V = 14.63, (B − V ) = +1.06. f The V and (B − V ) for LkHα-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). – 37 –

Table 5. Interstellar Absorption Line Propertiesa

Stars

A F B E D H LkHα- HD 101 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 +7 +9 +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

aVelocities are heliocentric, in km s−1; the equivalent widths are in mA.˚ bTwo minima in the NaI lines are at +7 and +16km s−1. The equivalent widths are with respect to the background of the stellar absorption line. cThe W s are with respect to the broad underlying stellar emission lines. dTwo minima in the NaI lines are at +5 and +13km s−1. – 38 –

Table 6. Radial Velocitiesa

−1 Feature Source v⊙ (km s ) Reference

Optical emission lines 1 +2 to +4 This paper Brα 2 +5 ± 2 Simon & Cassar (1984) Brγ 2 +6 ± 2 12CO (4.7 µm) absorption 10 +6.9 ± 0.5 Mitchell, et al. (1990) Hα emission 9 +2.9 ± 0.7 Fich, Treﬀers, & Dahl (1990) Peak of H92α emission 3 +10.3 ± 1.2 Knapp et al. (1976) 13CO(1-0)emission 4 +7:: 12CO (1-0) absorption 4 +7: H I (21 cm) emission 5 +8.7 Dewdney & Roger (1982) H I (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) a −1 At the position of LkHα-101, vLSR = v⊙ − 7.2km s .

Note. — Sources:

1. The HIRES spectra were extracted with a window 2′′ wide centered on LkHα-101.

2. Aperture diameter 3′′. 75 centered on the star.

3. Emission summed over an H II source 2′ in diameter centered on the star.

4. The 12CO and 13CO emission is from a region 2 - 3′ in diameter over which line structure is present between −3 and +11 km s−1. The narrow 12CO absorption at +7 km s−1 is seen against that background.

5. The velocity is an average over an extended region about 5′ 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 (1 - 2′) “obscuring cloud” in front of LkHα-101.

8. Average over a 45′′ beam centered on the star.

9. Hα emission averaged over a 2′ aperture centered about 1′ southwest of star.

10. Cold CO absorption against the stellar continuum at 4.7 µm. – 39 –

Table 7. Radio Source Correspondences

BW88 SO98 Table 1 Notes

8 1 ··· 5 ··· 2 ··· 5 9 3 66 4 5 4 ··· 2 ··· 5 ··· 1 ··· 6 ··· 1 ··· 7 194 LkHα-101 ··· 8 ··· 1 ··· 9 ··· 1 ··· 10 ··· 1 ··· 11 ··· 1 1 12 ··· 1 3 13 241 star D 2 14 ··· 1 7 15 ··· 1 6 16 48,49 3 4 ······ star E 2 ······ 5

Note. — 1: Nothing at radio position in either R or K 2: Nothing at radio position in R; outside area covered in K 3: There is a close double star (separation 1′′. 0) at the radio position; the brighter, southeast component has marginal Hα emission (W(Hα) ≈ 1.7 A)˚ 4: The star at the position has no grism- discernible Hα emission 5: Nothing at radio position in K; outside area covered in R – 40 –

Fig. 1.— A false-color image of NGC 1579 constructed from 300 s B, V and R-band exposures. The frame is roughly 7.′5 (1.5 pc for a distance of 700 pc) on a side, with north up and east to the left. – 41 –

Fig. 2.— The internal photometric errors of the optical and infrared photometry. – 42 –

303

233 187 30

192 100 107 304 D 205 44 132 63

95 72

151 78 157 194 27 111 215 225 140

83 122 105

126 118 253

139 70 112 243 180 32

Fig. 3.— A 300 s I-band image of NGC 1579 with the Hα-emission stars identiﬁed by their Table 1 numbers. LkHα-101 is star 194. – 43 –

Fig. 4.— A sample of GMOS spectra of Hα emitters. Each spectrum is identiﬁed by its Table 1 number. Prominent emission lines are labeled, along with Li I λ6707 absorption. Also indicated are some of the primary features utilized for spectral classiﬁcation. – 44 –

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 = 3.5 mag. Hα 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 LkHα-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 the text. – 45 –

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 LkHα-101 point has not been corrected for extinction. – 46 –

Fig. 7.— (J −H)vs. (H −K) color-color diagram for stars in NGC 1579. The symbols are the same as in Figure 5. Green points correspond to bright stars in Table 4 (including LkHα-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. – 47 –

Fig. 8.— The HIRES spectrum of LkHα-101 (on 2000 February 2) in the 6310 - 6395 A˚ region. – 48 –

Fig. 9.— The Fe II emission lines λλ6247,6249 in LkHα-101 on the three dates indicated, showing the variation in line structure near the peaks. The position of the whole line did not change signiﬁcantly. – 49 –

Fig. 10.— The spectrum of LkHα-101 between 6110 and 6140 A,˚ an average of 3 HIRES spectrograms. The wavelength scale is not corrected for the star’s velocity. Lines of the Mn II multiplet RMT 13 are indicated. – 50 –

Fig. 11.— The region of the Hα line in LkHα-101 on 2002 December 16. The section 6575 - 6620 A˚ has been expanded to show that the longward wing of the Hα line extends to at least 6600 A,˚ or a displacement of 1700 km s−1. Note that the [N II] λ6583 line is signiﬁcantly narrower than the Fe II; all the forbidden lines in this spectrum are single and narrower than the double permitted lines of Fe II, Ni II, and Mn II. Note the diﬀuse interstellar band at 6613 A.˚ Hα is shown at larger scale in Figure 12. – 51 –

Fig. 12.— The Hα and Hβ lines of LkHα-101 (on 2002 December 16). Note the diﬀerent intensity scales. The vertical line marks the mean velocity given by the narrow emission lines (+3.6 km s−1). – 52 –

Fig. 13.— The spectrum of LkHα-101 (above) and a bright section of NGC 1579 (below), obtained with the HARIS spectrograph. The intensity scales as plotted do not represent the true relative brightnesses of the two sources. – 53 –

Fig. 14.— Above: the 6540 - 6620 A˚ region in Star D (on 2002 December 16), showing the unusual structure of Hα (and the strong DIB at 6613 A).˚ Below: the spectrum of the G9 III-IV star HR −1 1787, here spun up to veq sin i = 80km s , to match approximately the line widths of star D. – 54 –

-40-20 0 20 40 1 1

0.8 0.8

0.6 0.6

0.4 0.4

0.2 0.2

0 0

-40-20 0 20 40

Fig. 15.— The predicted profile of a narrow emission line formed in a flat rotating annulus (of the dimensions of the horseshoe-shaped structure observed by Tuthill et al. 2002). As explained in the text, this particular calculation assumes that the annulus extends from 12 to 22 AU, is in Keplerian ◦ rotation about a 15 M⊙ star, and that its normal is inclined 30 to the line of sight. If the double Fe II lines observed in the optical region originate in such a structure, this is their expected profile, following blurring by the 7 km s−1 instrumental resolution. The ordinate scale has no significance. – 55 –

Fig. 16.— Left: the central region of NGC 1579, centered on LkHα-101, in the K-band; right: the same region in R. The area shown is about 2.′75 on a side, with north above and east to the left. – 56 –

To the dark observer lane

LkHα 101

illuminated surface of NGC 1579

bar, seen only at K

Fig. 17.— A schematic arrangement of star and nearby clouds intended to explain the diﬀerence (seen in Fig. 16) between the appearance of the LkHα-101 region in R and in the K-band. It is proposed that the star is heavily screened from the observer by the dark lane that crosses NGC 1579 from northeast to southwest (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 fraction of the star’s light passes through that dusty slab to illuminate the “bar”, and another fraction is scattered back through the slab toward the observer, but since extinction in the slab is wavelength- dependent, the “bar” will be seen preferentially at longer wavelengths. The weight of the various arrows suggests the relative amounts of energy passing along those paths. – 57 –

A. Table of Emission Lines in LkHα-101

This Table lists the stronger emission lines measured on the HIRES spectrograms of LkHα-101. The arrangement is as follows.

Col 1: the intensity-weighted mean wavelength of the line, not the wavelength of the intensity peak, and not corrected for the stellar velocity. This mean wavelength will change as the line structure varies. Many entries are followed by one of these characters:

d: double peaks

t: triple peaks

br: the line is broad

bl: the line is blended with another

as: the line is asymmetric

1: the line measured on only one spectrogram, usually because it falls in an interorder gap on the others

∗: a note follows at the end of the Table

Col. 2: W, the equivalent width of the entire line, in Angstroms. Lines of W less than about 0.05 A˚ are of low weight.

Cols. 3, 4: the likely identiﬁcation (with RMT number in parentheses) and its laboratory wavelength. A + means that there is probably another contributor. The [Fe II] and Fe II wavelengths are from Johansson (1977) and Johansson (1978), respectively.

Some lines of interest such as He I λ5875 do not appear because they fall in interorder gaps on all three spectrograms. – 58 –

Table 8. Emission Lines in LkHα-101

Mean Mean Ion Laboratory Mean Mean Ion Laboratory λ W(A)˚ λ λ W(A)˚ λ

4340.47 9. Hγ 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. Hβ 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 ? ··· – 59 –

Table 8—Continued

Mean Mean Ion Laboratory Mean Mean Ion Laboratory λ W(A)˚ λ λ W(A)˚ λ

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 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 ? ··· – 60 –

Table 8—Continued

Mean Mean Ion Laboratory Mean Mean Ion Laboratory λ W(A)˚ λ λ W(A)˚ λ

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 * Hα 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 ? ···

Note. — 5264: the line has a shortward wing, probably Fe II 5264.176 5746: the measurements include a shortward wing, possibly Fe II 5746.578 6456: the measurements include a distinct subsidiary peak on the shortward edge at about 6455.69, probably a blend of 2 weaker Fe II lines, plus a contribution from O I RMT 9 6548: The [N II] line is confused by an overlapping atmospheric H2O line 6562: see the text for a discussion of Hα 6678: the line falls in a corner of the frame, and possibly for that reason measurements are discordant