HST/NICMOS Observations of Variability in the Scattered Light from Young Stellar Objects

Angela S. Cotera

SETI Institute

[email protected]

Glenn Schneider

Steward Observatory

[email protected]

Dean C. Hines, Barbara A. Whitney

Space Science Institute

[email protected],[email protected]

Karl R. Stapelfeldt

NASA Jet Propulsion Laboratory

[email protected]

Deborah L. Padgett

Spitzer Science Center

[email protected]

ABSTRACT

We have obtained a second epoch of Hubble Space Telescope NICMOS observations of seven Young Stellar Objects: HH 30, L1551 IRS5, IRAS 04016+2610, IRAS 04302+2247, Haro 6-5B, CoKu Tau/1 and DG Tau B. The first epoch observations were taken in 1997 as part of two different programs, while the second epoch were obtained in 2004 as part program to measure polarization in these objects. Data obtained during the first epoch utilized a variety of NICMOS filters, the results of which have been previously presented in the literature, with the exception of L1551 IRS5. Therefore, the multi-filter data obtained during the first epoch for L1551 IRS5 is also presented here. Data obtained as part of the second epoch included the polarization filters and the broad band F160W filter. All of the objects show significant variability in the scattered –2–

light patterns observed at 1.6 µm between the two epochs. For three of the objects, the variability is most likely associated with binary motion, while the remaining objects all exhibit variations consistent with a rotating photospheric hot spot. Additional consideration is given to the production of the observed changes as a result of variable extinction located close to the protostar.

Subject headings: stars: formation

1. Introduction

Sixty years ago, T Tauri stars (TTS) were first distinguished as a distinct class of objects based on their photometric and spectroscopic variability (Joy 1945). In subsequent years, TTS were further sub-typed into weak-lined T Tauri stars (WTTS) and classical T Tauri stars (CTTS) in reference to the strength of observed Hα emission (e.g. Bertout 1989). As the evolutionary status of TTS was related to other prestellar objects, they have become known more generally as Young Stellar Objects (YSOs). Based on observations in the ρ Ophiuchus region, Lada & Wilking (1984) suggested that YSOs be sub-classified into three categories based on the near-infrared (NIR) excess – Class I to Class III – which approximates an evolutionary sequence from an early protostellar to the more typical T Tauri stars (Adams, Lada & Shu 1987). An earlier stage – Class 0 – was first suggested by André et al. (1993), while evidence of debris disks around more evolved stars (e.g. Backman & Paresce 1993) has established a bridge from YSO studies to main sequence stars. This evolutionary scheme primarily delineates the disappearance of the protostellar envelope, the formation of a circumstellar disk, and the eventual dissipation of that disk.

The inherent variability of YSOs is observed as changes in the photometric brightness, the spectroscopic equivalent width of emission and absorption lines, and the high resolution morphology of the light scattered by material in the circumstellar environment. The photometric variability has been long established (Joy 1945), with both periodic and aperiodic light curves observed at both optical and near-infrared wavelengths (NIR, e.g. Basri & Batalha 1990; Skrutskie et al. 1996). Spectroscopically, the observed “veiling” (so called because the absorption lines are not as deep as comparable standard stars) has been attributed to accretion onto the stellar photosphere (e.g. Basri & Batalha 1990), with the variability related to changing accretion rates. The timescales of the observed variability are not uniform however, and range from days to years. Periodic photometric variability has been observed in many TTS, although for many known variable objects possible periodicity has not been constrained as only a few observational epochs are available; other TTS stars exhibit purely irregular variability (Bouvier et al. 1995).

Ground-based observations of the variability in both the measured brightness and the spectroscopic veiling suggests that the physical causes arise from one or both of two different fundamental phenomena: (1) temperature variations of the stellar photosphere, and (2) uneven illumination caused by a non-uniform circumstellar environment. Herbst et al. (1994) proposed a classification –3– scheme for the variability suggesting that Type I variations were the result of the rotation of cool spots and associated primarily with WTTS. Type II variations are the result of a combination of hot spots and variable accretion and are associated predominately with CTTS. Finally, Type III variations results from variable extinction by circumstellar dust and are primarily associated with UXors, which Herbst et al. (1994) denote early T Tauri stars (ETTS), and are often associated with Herbig AeBe stars. There is not a one-to-one correlation between TTS classification and variability type, but the proposed scheme does provide a coherent framework for discussion of YSO variability.

The Taurus-Auriga cloud contains some of the best studied YSOs, often invoked as prototypes for low mass, solar type, star formation. Although there is now evidence that the Sun formed in an environment more like Orion than Taurus-Auriga (Hester & Desch 2005), at a distance of ∼140 pc, high resolution observations enable us to study Taurus YSO structures on spatial scales of tens of AU. For example, early serendipitous Hubble Space Telescope (HST) observations of the Class II object HH 30 (e.g. Burrows et al. 1996) provided an almost textbook example of theoretical models (e.g. Whitney & Hartmann 1992) of light scattered from a nearly edge-on (i∼82o), ﬂared, optically thick YSO accretion disk with a diameter of about 450 AU. The observed variability of the scattered light nebulosity associated with HH 30 (Burrows et al. 1996; Stapelfeldt et al. 1999; Watson et al. 2006) is not surprising, given the ubiquitous nature of TTS variability. Wood & Whitney (1998), investigating the observational signatures of stellar hot spots, suggested that the small variations in observations HH 30 by Burrows et al. (1996) could be explained with such a model. Their model invoked a hot spot created via magnetic stellar accretion mediated by a dipole ﬁeld not aligned with the stellar rotation axis. Wood & Whitney (1998) predicted that the disk should brighten asymmetrically during a stellar rotation period; illuminating the disk like a lighthouse beam sweeping over the ocean. A major lateral shift in the brightness distribution within the HH 30 disk was subsequently reported by Stapelfeldt et al. (1999), however a connection between the stellar rotation period and the observed variability remains elusive (Watson et al. 2006).

As part of HST program GO 10178 (PI: Hines), which utilizes polarimetry with both NICMOS and ACS to investigate the dust evolution in the disk around YSOs to be presented in a future paper, we obtained NICMOS F160W observations in 2004 (Cycle 13) of seven objects which had previously been observed in the same bandpass with NICMOS in 1997 (Cycle 7). Two of the objects, L1551 IRS 5 and HH 30 were observed as part of GTO 7228 (PI: Young), while the other five, CoKu Tau/1, DG Tau B, Haro 6-5B, IRAS 04016+2610 and IRAS 04302+2247 were observed as part of GO 7418 (PI: Padgett). All of the objects show significant variability in their scattered light patterns between the two observational epochs. In this paper we present a comparison of the F160W observations taken at the two epochs, and discuss the observed variations. In section 2, we present the observations and data reduction for both epochs, in section 3 we discuss the individual objects, in section 4 we discuss the implications of the observed variations in their ability to discriminate between the proposed mechanism, and in section 5 we summarize our findings. –4–

2. Observations and Data Reduction

All observations were taken with the Hubble Space Telescope, using the NICMOS camera. As can be seen in the observation log (Table 1), the integration times were different for each of the different HST programs. The images were reduced and calibrated with procedures developed by the NICMOS IDT (Schneider 2002). The raw MULTIACCUM images were calibrated in an IDL-based analog to the STSDAS CALNICA task. “Synthetic” dark frames, high S/N calibration reference flats, and linearity files appropriate for HST Cycle 13 were supplied by STScI. Dark frames for the Cycle 7 data were obtained from observed, on-orbit darks. Instrumental count rates were converted to physical flux densities based upon absolute photometric calibrations established in the HST SM3B program (see below). After conversion to flux units, the images were distortion corrected by mapping input pixels with X:Y scales of (75.950, 75.430) mas/pixel onto a rectilinear grid of 75.95 mas square pixels using the IDL program IDP3, developed by the NICMOS IDT (Schneider & Stobie 2002). All data were subsequently shifted and combined using IDP3. The science objective for the Cycle 13 observations was to measure the polarization of the target objects, therefore the only broad band data obtained in 2004 was a reference image taken with the F160W (1.6 µm) filter. Consequently, only the F160W data from the previous programs has been re-reduced and presented here. An observing log, with the dates and integration times for the two epochs of observations, is presented in Table 1.

The first epoch observations of (Padgett et al. 1999), have higher signal-noise due to the significantly longer integration times. The second epoch observations were made after the installation of the NICMOS cryo-cooler in SM3B, and benefit from improved detector quantum efficiency due to the increase in the operating temperature from 62 K to 77.1 K. The change in quantum efficiency is reflected in the conversion factors from ADU/s to Jy of 1.498×10−6 for Cycle 13 and 2.061×10−6 for Cycle 7 in the F160W filter1, an increase in sensitivity of approximately 1.38. Padgett et al. (1999) did not use the photometric calibration provided by STScI since the pipeline calibration assumes a constant flux per unit wavelength, and the SEDs of their program sources were significantly different. They compared model and previously observed SEDs for their objects and derived ADU/s to Jy calibrations for each object in each filter. They found that their derived flux values differed from those based on the photometric calibration provided by STScI by 25%-180%. Since our primary interest here is to compare the morphological differences between the two epochs of observations, we use only the standard calibration. This provides robust results, unless there is an unknown significant color change over the bandpass between the two epochs. The derived aperture fluxes presented in Table 1 use the same radius as Padgett et al. (1999), therefore differences in flux values between this paper and Padgett et al. are essentially due to their use of customized calibration values (although the general calibration has also changed slightly since their data were published). Despite these caveats, however, it is clear that there are significant differences between the two epochs.

1http://www.stsci.edu/hst/nicmos/performance/photometry, March 2006 –5–

After flux calibration and pixel rectification, each object was rotated to celestial coordinates using the bi-cubic interpolation routines contained within the IDP3 software (Schneider & Stobie 2002). In order to align the observations taken at the different epochs, we first attempted to use only upon the astrometric information provided with the images. The accuracy of alignment using this method, however, is dependent upon numerous factors, primarily the accuracy of the coordinates in the HST Guide Star Catalog for both the dominant guide star and the roll guide star. Utilizing high accuracy astrometric catalogs to revise the guide star coordinates still results in absolute positional uncertainties of ∼70-80 mas (see §3.2) which is ∼1 pixel. Fortunately, for our objects (with the exception of HH 30 and IRAS 04302+2247) at 1.6 µm there is a strongly peaked source associated with the YSO which we are able to use for an accurate centroid determination, which allows for alignments to within ∼0.1 pixel. We therefore registered the two epochs based on alignment of the centroids of highly peaked sources. For HH 30 and IRAS 04302+2247, we used a combination of point or highly peaked sources in the frame, and minimization of residuals in the subtracted images. Therefore, the method we used to obtain all the differenced images presented here does not take into account proper motions of the primary source.

Photometry results are presented in Table 1. We performed aperture photometry utilizing 50 pixel (∼3.6) radius apertures, and subtracting the background based on regions without emission on the NICMOS images. In all cases, the uncertainties in the total flux within the aperture is 0.1 mJy. In addition, except for L1551 IRS 5, the objects have a distinct bipolar morphology, consistent with models of starlight scattered off nearly edge-on flared optically thick disks and surrounded by circumstellar envelopes. We have used IDP3 to measure the flux and background within irregular polygons which outline these individual bipolar nebular lobes. In Table 1 and all subsequent discussion, we use “A” to indicate the brighter of the two nebular lobes, and “B” when discussing the fainter lobe. The nomenclature used in Table 1 and throughout the paper is illustrated in contour plots of the regions shown in Fig. 1. The uncertainties for these measurements are 5%.

3. Results

In addition to the sometimes large variability in the overall flux levels, all of the images show significant changes in the spatial distribution of their observed scattered light patterns. Grayscale images from each epoch, and a differenced image (with the first epoch subtracted from the second epoch), are presented in Figures 2, 8, and 10-14. As can be seen in Table 1, there have been significant photometric variations between the two epochs. Therefore, in order to enable us to focus on the morphological changes, we have normalized the images to the flux in the first epoch, using the ratio of the large aperture fluxes from Table 1. The subtleties in the variations, however, are best seen by blinking the aligned image frames at a variety of display ranges; short movies –6– showing the blinked images are available online2. Although many of the variations discussed below are perhaps more readily apparent in the blinked images, we refer primarily to the figures presented in the paper in the following discussion of the individual objects.

3.1. HH 30

Whitney & Hartmann (1992) predicted that the light scattered off flared circumstellar disks surrounding Class I–II protostars could be imaged in systems where the central star is obscured by an edge-on disk. The first example of a disk detected in this manner was that of HH 30 by Burrows et al. (1996). Their WFPC2 images resolved the HH 30 nebulosity into a characteristic nebular pattern for an edge-on disk: two oppositely concave nebulae separated by a dark dust lane (Whitney & Hartmann 1992), creating a nebula with bipolar symmetry. The HST WFPC2 observations of HH 30 also showed that the scattered light pattern is highly variable, having changed considerably in the brightness pattern observed between the two epochs presented by Burrows et al. (1996). They found that one side of the scattered light nebula faded while the opposite side brightened. Wood & Whitney (1998) suggested that the observed asymmetry in HH 30 could be explained by a single hot spot created via magnetic stellar accretion mediated by a dipole field not aligned with the stellar rotation axis. Based on additional epochs of HST observations, Stapelfeldt et al. (1999) observed the asymmetry predicted by Wood & Whitney (1998), and constrained the timescale of the variability to be <3 years. Stapelfeldt et al. (1999) found that this timescale precludes changes in the outer disk as an explanation, and established that the cause of the notable time variable lateral asymmetry must be an illumination effect originating in the inner few AU of the disk. They found that the limited data was consistent with both the hot-spot models and possible shadowing by vertical disturbances in the inner disk, which they speculated might be caused by bars, warps, or spiral arms. Recently, O’Sullivan, et al. (2005) successfully reproduced the ground-based observed variability of AA Tau by vertical shadowing from a warped inner disk, but similar modeling of HH 30 is currently unavailable.

Over the past decade, HH 30 has been observed extensively both with HST at 16 epochs (Watson et al. 2006), and from the ground (e.g. Wood et al. 2002; White & Hillenbrand 2004). Wood et al. (2002) found that their ground based data was suggestive of an 11.6 day and 19.8 day periodicities, which is consistent with a stellar origin for the observed variability. Multi-epoch HST observations of the scattered light lateral asymmetry (e.g. Fig. 2), have not conﬁrmed this periodicity; however, the temporal coverage has not provided conclusive constraints (Watson et al. 2006). If the lateral asymmetry is produced by hot spots caused by magnetic accretion, any periodicity should be tied to the stellar rotation and thus relatively short (e.g. 18 days Wood et al. 2002; Watson et al. 2006), and might be expected to have strong veiling due to variable accretion. To date the observations have not been able to provide a deﬁnitive answer to which mechanism is

2http://nicmosis.as.arizona.edu/∼cotera/variability –7– responsible for the variability in HH 30, and indeed magnetic accretion hot spots and circumstellar obscuration may both be at work.

Our observations of HH 30, which provide data only for two epochs of NICMOS observations are shown in Fig. 2. We see the same variation in the lateral asymmetry as seen in Burrows et al. (1996), Stapelfeldt et al. (1999), and Watson et al. (2006). In our data, we see that in the brighter nebula (lobe A in Fig. 1), the peak emission has shifted from the left in first epoch to therightinthesecondepoch.Atthesametime,theemissioninthefainternebula,lobeB,has shifted from the right to the left (although this is not readily apparent in Fig. 2). This is the same pattern of variability observed and modeled previously. Watson et al. (2006) were unable to reach definitive conclusion as to the mechanism responsible for the variations seen in their 16 epochs of data. The two epoch, single filter observations presented here cannot provide additional insight considering the extensive previous research done on HH 30; however, the data provides valuable reference information when examining the observed variations in our remaining sources. We note, however, that despite the clear variations in the scattered light patterns, the overall magnitude variation in HH 30 is the lowest of those objects presented here (see Table 1), with ∆MF 160W <0.1.

3.2. L1551 IRS 5

L1551 IRS 5 is, in fact, a binary system. High resolution (100 mas) 2 cm VLA data (Bieging & Cohen 1985) first suggested that IRS 5 consisted of two stars rather than one. Rodríguez et al. (1986), repeating the VLA observations at 2 cm, also resolved two sources, but attributed the emission to the edges of a shock-ionized torus. At millimeter wavelengths, 2.7 mm observations of Looney, Mundy & Welch (1997), and 7 mm observations by Rodríguez et al. (1998), both resolve two sources, with the same separation and position angle as the 2 cm data, and ascribe the emission at these wavelengths to hot dust immediately around the stars. The case for the binarity has been further bolstered by the detection of two distinct jets at optical (Fridlund & Liseau 1998), near- infrared (Itoh et al. 2000), and radio (Rodríguez et al. 2003b) wavelengths.

3.2.1. First Epoch Observations

In 1997 (HST Cycle 7), observations of L1551 IRS 5 were obtained in the F110W, F160W, F187N, F204M, and F212N ﬁlters. This has not been previously published, so we present the results of those observations here. In Fig. 3 we present a color composite image of the F110W, F160W and F204M data. In Fig. 4, we present a contour map of the 1997 F160W image, in which we label the observed morphological features.

In Fig. 4, the contours delineating the northern jet observed by Fridlund & Liseau (1998) and Itoh et al. (2000) are readily apparent as is region D (nomenclature Neckel & Staude 1987), consistent with previous work indicating that the northern jet terminates at region D (Fridlund & –8–

Liseau 1998). Due to the short integration time, the second (southern) jet is not obvious. Near the north northwest edge of the reflection nebula, there is a small, but real “clump” of excess flux at all of the wavelengths which can be seen in Fig. 3. In Fig. 4, we have also plotted as a first approximation, a parabola surrounding the clump to emphasize what we describe as a “wake”. The morphology of these two features strongly suggests they are associated with the scattered light from IRS 5. The clump feature is separated from the peak in IRS 5 by approximately 1.6 (∼225 AU), which is below the resolution of most published images of the dust and gas, so a resolved counterpart at other wavelengths would not be expected from current data. The fact that it does not appear on any of the high resolution centimeter or millimeter images of the region, however, suggests that the clump is likely the result of light scattered off a dust feature, rather than an indication of an additional protostar.

In the NICMOS images, the reflection nebula clearly shows the characteristic concave shape produced by scattering off an optically thick flared circumstellar dust disk. If we simply use the contours obviously associated with the northern jet and region D to draw a line defining a possible jet axis, and trace this back to a possible origin (see Fig. 4), we find the line passes through the apex of a parabola fitted to the ridge of peak emission in the reflection nebula. Using this geometry, we would conclude that the protostellar source is located behind the apex of the parabola, which has been the assumed location in most previous work (e.g. Lucas & Roche 1996; Itoh et al. 2000). In addition to the concave reflection nebula, however, there is also a very strong, circularly symmetric peak in the emission observed in the NICMOS images, located south east of the apex.

The compactness and circular symmetry of the observed NIR peak, which becomes more point like as we move to longer wavelengths, strongly suggests that this peak may in fact be the location of at least one of the two protostars in the region. In order to investigate the nature of this compact source, we first need to establish how strongly peaked the emission is, and whether it is a point source or a resolved compact source. In Fig. 5, we present a line cut through the F110W, F160W and F204M images extending from the north northeast to south southwest, along the southern ridge of the nebula and through the compact source. The peak flux from the compact source varies from ∼1.4 – 2.9 times that of the surrounding scattered light nebula as we go from the F110W to the F204M filters (1.10 – 2.04 µm respectively). We created an artificial point-spread function for the F110W, F160W and F204M filter images using TinyTIM (Krist & Hook 1997), with the parameters adjusted to reproduce the conditions of the telescope at the time of our observations. By scaling these model PSFs so the peaks match the observed data, and overplotting them onto these line cuts (Fig. 5), we see clearly that although the observed flux is strongly peaked, it is extended beyond that expected for a point source at all wavelengths. The object is therefore not a point-source but an extended source, consistent with the classification of IRS 5 as a still embedded Class I YSO.

The question then becomes, how does this NIR peak relate to the known radio and millimeter sources whose coordinates are presented in Table 2? There are no reference NIR/radio objects within the field of view, therefore the only possibility for an unbiased determination of the position –9– of the NIR peak is to utilize astrometry provided by HST. The generally quoted uncertainty for HST astrometry is ∼1, which is insufficient to accurately compare the position of our NIR source to the high resolution images at radio and submillimeter wavelengths with uncertainties of 140 mas (see Table 2). Upon investigation, we determined that our relatively large positional error had been introduced by inaccurate coordinates for the HST guide star, HD 285845, used to establish the image coordinates. We therefore used more accurate positions for the reference star provided by the ACT Reference Catalog (Urban et al. 1997) to revise the coordinates of our guide star, resulting in an applied offset of ∆RA= 0.035s in ∆Dec=0.014, with positional uncertainties of 35.3 mas in RA and 29.8 mas in Dec. Coordinates for the roll angle guide star were also corrected with more recent values (GSC-ACT catalog; Lasker, Russel & Jenkner 1996); the newer values necessitated a slight increase in the roll angle (0.◦0274), with an additional positional uncertainty at the location of the NIC2 array of 9 mas. The total uncertainty in the aperture position is ∼100 mas (M. Lallo 2001, private communication), which we take to be equivalent to 70 mas in both RA and Dec. We therefore applied both a translational correction to the image header coordinates based on the above offsets, rotated by the corrected roll angle, and propagated the uncertainties introduced from the guide star positional error, roll angle, and aperture position in quadrature. Our derived position for the NIR point source is presented in Table 2; the positional uncertainties are σRA=80 mas and σDec=77 mas. When all of the positions in Table 2 were initially plotted on our astrometrically corrected F160W contour map, the correspondence between the NIR features and previous measurements showed significant scatter. Even when the most accurately measured source positions available at the time, the radio and millimeter observations of the binary sources, were plotted (see Fig. 6), there is still ambiguity between the compact NIR source and the observed binary sources. Upon closer examination, we noticed that the published positions did not agree with each other to within their quoted uncertainties. The two VLA observations from the mid-1980’s (Bieging & Cohen 1985; Rodríguez et al. 1986, BC85 and R86 in Table 2) seemed to agree well, as did the observations from the late 1990’s (Looney et al. 1997; Rodríguez et al. 1998, L97 and R98 in Table 2).

A study of pre-main sequence stars in Taurus-Auriga by Frink et al. (1997), derived proper motions for 72 T Tauri stars located in the central regions of the cloud. The quoted proper motions are on the order of ∼3-150 mas/yr, with µα ∼ 12±4.3 mas/yr, and µδ ∼−20±4.3 mas/yr given for L1551-51, one of the closest objects to IRS 5 in Frink et al. (1997). Therefore, when we first analyzed the data, we strongly suspected that our inability to properly align the various reference sources resulted from the proper motion of L1551 IRS 5. Using the coordinates of the high resolution 1980 and 1990 observations (Bieging & Cohen 1985; Rodríguez et al. 1986; Looney et al. 1997; Rodríguez et al. 1998), differencing the results in all possible permutations and propagating the uncertainties, we derived a proper motion of µα ∼ 8.9 ± 5.3 mas/yr and µδ ∼−21.3 ± 8.1 mas/yr; which was in good agreement with the proper motion of L1551-51. The large uncertainties for our estimate are the result of unexplained differences in the 1983 and 1985 VLA 2 cm positions. Our results, based on the literature at the time of the initial observations, is in good agreement with the more recent –10– results of Rodríguez et al. (2003a). They reanalyzed their radio data from 1983, 1995 and 1997, and derived proper motions of: IRS 5A (North) µα ∼ 11.7 ± 1.9 mas/yr, µδ ∼−20.3 ± 2.5 mas/yr, and IRS 5B (South) µα ∼ 14.7 ± 1.6 mas/yr, µδ ∼−22.2 ± 2.8 mas/yr. With the improved absolute astrometry of the HST images, and applying our derived proper motion to the previous radio and millimeter observations, we find that the position of the northern binary source, IRS 5A, unambiguously corresponds with the strong compact NIR feature seen in the NICMOS images of L1551 IRS 5 to within 2σ positional uncertainty (see Fig. 2). There is no NIR feature which corresponds to the binary companion IRS 5B. Given the alignment of features within the system, we suggest that the observed NIR light results primarily from emission by IRS 5A scattered off a flared circumstellar disk, with virtually no indication of additional emission, scattering or absorption from IRS 5B.

To investigate whether the suggestion that the scattered light is associated only with IRS 5A is reasonable, we performed an analysis of the local extinction. We created an extinction map by first convolving the F160W image with a gaussian function, degrading the image resolution to match that of the F204M image. We then converted the images to magnitudes and subtracted them to make a F160W–F204M (approximately equivalent to H–K) color map, and converted the colors to an equivalent AV using the extinction law of Rieke & Lebofsky (1985, Fig. 7). In the extinction map, we find that the compact source is slightly, but not significantly more extincted than the surrounding nebulosity. The dust clump discussed above, however, does appear to be near a local extinction maximum. The wake is also slightly reddened, suggesting either more light has been preferentially scattered at longer wavelengths into these regions, or there is greater extinction, or both. The most striking aspect of the color map, however, is the large increase in extinction to the northwest of the compact NIR source, where the peak extinction reaches AV ∼30. The peak in the extinction likely indicates the location of a circumstellar disk.

The extinction for IRS 5 has been estimated to be AV ∼19 (e.g. Carr et al. 1987), and from our observations we derive a peak extinction of AV ∼30. These estimates, however, are based on scattered light and do not represent the true extinction in the disk. Based on ISO observations, White et al. (2000) derive AV ∼130 in the disk. Radiative scattering models of HH 30, a significantly more evolved system, indicate that in this nearly edge on system, the extinction from the disk may be up to AV ∼ 7, 900 (Cotera et al. 2001). Therefore at NIR wavelengths, IRS 5B could be completed extincted by the material surrounding IRS 5A if it is currently behind IRS 5A: but we might still expect to see light scattered off a circumstellar disk surrounding IRS 5B. If the IRS 5B star+disk system were in front of IRS 5A, on the other hand, we would expect to at least see some absorption of the scattered light from IRS 5A. Based on our astrometry, the observed jet is most consistent with an origin in IRS 5A, in disagreement with Rodríguez et al. (2003b). The proposed juxtaposition of IRS 5A and IRS 5B is also satisfying in that it explains why we do not see strong evidence of a second jet in our images, since near the star the base of the jet would also be extincted at NIR wavelengths. Testing whether our proposed geometry for this complex region is correct, must await the results from our detailed radiative transfer models to be presented with –11– our analysis of the polarimetry data taken as part of our second epoch observations.

3.2.2. Comparison with Second Epoch Data

To compare our second epoch of data to the first epoch, we have used the IRS 5A NIR peak for alignment, thereby neglecting any differences associated with the proper motion of the object during the 7 years between observations. The two epochs, and the differenced image are presented in Fig. 8. We note that all of the the features discussed above in association with first epoch data, are seen in the second epoch observations at higher signal-to-noise. The first striking difference between the two epochs is the fact that we can see the propagation of the jet we have associated with IRS 5A in the 1997 observations, outward to the southeast. The direction of the most highly collimated jet outflow, nearest to the star appears to have moved significantly to the north. Our first consideration was whether the rotation information provided by HST astrometry was incorrect, but careful investigation of the primary guide star and the roll guide star coordinates indicated an error no greater than that found in the first epoch observations (∼0.◦04). Alignment of the base of the northern jet between the first and second epochs requires an additional image rotation of ∼3. ◦8 (see below). Therefore we conclude that the observed change in the orientation near the base of the jet is most likely a physical effect. The observed movement would be consistent with the suggested precession in the jet identified in previous observations as a “wiggle” in the jet (Itoh et al. 2000; Hartigan et al. 2000).

We find that if we define a reference jet axis as the line between the NIR peak and the brightest knot in the first epoch of data, this knot propagates linearly outward. If we trace the peak of the jet at both epochs, we find that the knot is actually an inflection point in a non-linear jet outflow. This is illustrated in Fig. 9, where we present contour plots of the jets from both epochs, overlaid with lines which trace the peak of the jet emission in both epochs. When we align the inflection knot, as shown in the lower right panel of Fig. 9, the observed differences are consistent with what can best be described as circular motion about the inflection point of the jet. (This subtle motion is best seen in the blinked images. 3) The blinked images suggest that in addition to linear outflow, the jet also contains a transverse component, with the inflection point indicating where the transverse component changes sign.

The bright knot which defines the inflection point, has moved a projected distance of ∼95 AU, assuming a distance to L1551 IRS 5 of 150 AU (Kenyon et al. 1994). Using the dates of the two observations (see Table 1), this results in a projected velocity of ∼65 km s−1. If we assume an outflow angle of 45o − 60o, and that the emission is blueshifted (e.g. Fridlund et al. 2005), this would mean an actual velocity of −92 to −130 km s−1. This is in reasonable agreement with the spectroscopic results for the F3 emission knot of Fridlund et al. (2005). Finally, in the differenced

3http://nicmosis.as.arizona.edu/∼cotera/variability –12– image (Fig. 8), we ﬁnd a very bright knot of emission near the base of the IRS 5A jet in the second epoch data, located approximately 175 AU from the NIR peak.

When we further examine the differenced image, we find that the variation between the two epochs does not resemble at all the variable bilateral asymmetry seen in HH 30. In fact, the observed differences between the two epochs are best explained by a rotation of the scattered light pattern. In particular, if we rotate about the NIR peak until the axes defined by the peak and the nearest, well collimated jet emission are aligned, we find that the major features seen in the scattered light patterns are also nearly aligned. Using this method we find that the angle of rotation is ∼3.8o.We note that of all the objects presented here, only L1551 IRS 5 requires rotation of the image about the compact source used to align the main observed features. When combined with our estimate of the small uncertainties introduced by calculation of the roll angle by the spacecraft, we find that the most likely explanation is we are seeing real effects.

Rodríguez et al. (2003a) estimate a binary rotation period for the system of 260 yr. Assuming this period, for a face-on system we would expect to detect motion consistent with ∼9.8o between our two epochs. Assuming an inclination of 45o, our observed rotation angle is consistent with the suggested binary period. Interestingly, the peak of the bright “clump” of scattered light noted in first epoch, has also shifted. This supports the idea that we are observing a clump of dust which is scattering the light from IRS 5A, since at its projected distance from IRS 5A (400 AU), coordinated physical movement seems highly unlikely. The clump could be more extended than it appears in the images, such that as IRS 5A orbits, different portions of the clump are illuminated. The numerous questions raised by these observations and the proposed juxtaposition of sources within L1551 IRS 5 will be tested in our upcoming models which will include the observed polarization.

3.3. IRAS 04302+2247

Designated the “Butterfly star” by Lucas & Roche (1997) because of its distinct morphology (Fig. 10), IRAS 04302+2247 is an edge-on system, with a wide dark lane observable at NIR wavelengths. IRAS 04302+2247 has Lbol ∼0.3L (e.g. White & Hillenbrand 2004), and a calculated circumstellar mass from 1.3 mm observations of 0.04 M (within a radius of 4200 AU, Motte & André 2001). Duchêne et al. (2004) classify IRAS 04302+2247 as a flat spectrum source, because of its near-infrared to mid-infrared spectral index; and suggest that it is a transitional object between Class I and Class II. In general, however, IRAS 04302+2247 has been found to be a Class I YSO (e.g. Kenyon & Hartmann 1995; Lucas & Roche 1997; Padgett et al. 1999; White & Hillenbrand 2004), with the NIR emission consistent with scattering from both a disk and an envelope.

Padgett et al. (1999) found that the thickness of the dark lane observed with HST NIC- MOS decreases by 30% going from the F110W (1.10 µm) to F205W (2.05 µm) filters, providing strong evidence for the interpretation of the dark lane as the result of an optically thick circum- –13– stellar disk. IRAS 04302+2247 also has spatially variable extinction within the scattered light pattern, consistent with a circumstellar envelope. Wolf, Padgett, & Stapelfeldt (2003) modeled IRAS 04302+2247 as a disk+envelope system. They found that the observed NIR and submm morphology is consistent with light scattered off a flared circumstellar disk with an outer radius of 300 AU, and a scale height h(100)=15 AU. Wolf et al. (2003) also found that the dust grain distribution is consistent with grain growth similar to HH 30 (Cotera et al. 2001).

Lucas & Roche (1997) were the first to note the interesting quadrupole symmetry, labeling the quadrants P1-P4, corresponding to the northwest, southwest, southeast and northeast lobes. In order to be consistent with our own nomenclature, we refer to the regions as A1 and A2 for the north and south western quadrants, B1 and B2 for the north and south eastern quadrants, respectively (see Fig. 1). We have used a polygon aperture to measure the flux by quadrant and present the results in Table 3. Between the first and second epochs, there has been an overall decrease in the emission of ∼26%. The western lobe (A, see Table 1) has decreased by ∼35% while the eastern lobe (B) has decreased by only ∼15%. Lucas & Roche (1997) also observed variations in the flux from the different lobes, even within observations taken on the same night. They found that the emission from B1 remained constant in all their observations, while A2 had an average overall change of ∼30% between their observations in January and November 1995. Surprisingly, we have very similar results, with B1 remaining constant to within the uncertainties, and A2 exhibiting a change on the order of 35%.

In order to make a more detailed comparison of our data with that of Lucas & Roche (1997), taken at lower resolution, we turn to the flux ratios between the different combinations of quadrants for the various observations. If the variability of IRAS 04302+2247 is caused by bipolar, symmetric hot spots, we would expect the ratios of the northern to southern lobes, on both sides of the disk, to be inversely proportional; ie. if the ratio of A1/A2 increases, the ratio of B1/B2 should decrease. We do not find strong evidence for this expected asymmetry, although for our observations the ratio of B1/B2 increased by 35%, the ratio of A1/A2 decreased very slightly but the difference was less than the measurement uncertainties. Re-evaluating the results of Lucas & Roche (1997), we find the average of their January and November 1995 observations are more consistent with the expected axisymmetry, with the ratio of B1/B2 (their P1/P2) increasing by 5%, while the average ratio of A1/A2 (their P4/P3) decreased by 37%.

In Fig. 10, we have normalized the flux levels between the two epochs prior to differencing the images using the ratio of the overall flux as presented in Table 1. This emphasizes the morphological differences in the brightness distribution. In the differenced image presented in Fig. 10, we see that the variability is approximately consistent with that seen in HH 30. In the normalized images, the B1 and A2 quadrants are generally brighter, while the B2 and A1 quadrants are dimmer. The differences with HH 30 arise from the fact that the observed brightness distribution is clumpy. In addition to the overall variations, there are also observed morphological variations within most of the quadrants, indicative of an underlying system of greater complexity than the nearly ideal situation in HH 30. –14–

3.4. IRAS 04016+2610

Also known as L1489 IRS, only one lobe of the typical bipolar nebula is immediately obvious in the scattered light images of IRAS 04016+2610 (Fig. 11); although the fainter outﬂow cavity is clearly seen in polarized intensity maps (Lucas & Roche 1997; Cotera et al. 2006). IRAS 04016+2610 is considered a Class I YSO (e.g. White & Hillenbrand 2004); however, based on submillimeter continuum observations Hogerheijde & Sandell (2000) suggest it may be an object in transition from a Class I to a Class II. IRAS 04016+2610 has Lbol ∼3.7 L (Kenyon & Hartmann 1995), and a mass of 0.03 M derived from 1.3 mm observations (Motte & Andr´e 2001). Hoger- 18 heijde et al. (1998) derived a mass of 0.013 M based on their observations of C O emission over 20×20. Based on the position-velocity observations in HCO+, the gas may be orbiting an object with 0.9 M (Hogerheijde et al. 1998).

In the HST NICMOS F205W image, Padgett et al. (1999) found that the observed ∼ 2 µm flux is dominated by the central point source. They also determined that there is a prominent dark lane oriented at a PA∼80o, consistent with the observed molecular gas, and perpendicular to the observed outflow at a PA∼165o (Hogerheijde et al. 1998). The location of the dust lane is illustrated in the contour plot (Fig. 1). The geometry is consistent with a system which has an infalling envelope, bipolar outflow cavities and a flared circumstellar disk at a inclination of ∼60o (Kenyon et al. 1993). The inclined disk is most likely occulting the majority of the scattered light from the northern lobe. The detailed geometry of IRAS 04016+2610 , however, is not as straight forward as other objects in this study. Although there is a definite dark lane located just north of the point source, in the broadband NIR images there in an additional, approximately linear, extinction feature parallel to the main dust lane ∼7 south of the primary dust lane (see Fig. 1). Padgett et al. (1999) noted this feature and suggested it is a result of localized extinction; however, the straight relatively long feature is not consistent with clumpy dust. The observed extinction most resembles absorption by a disk; but a disk at this location seems implausible.

Observations in the submillimeter continuum show a second peak of emission, northeast of the primary peak (Hogerheijde & Sandell 2000), consistent with the system containing a close binary (∼20 AU). In the NIR observations of the scattered light (Fig. 11), however, only one star is readily apparent. Wood et al. (2001) modeled IRAS 04016+2610 as a star+disk+envelope system, with two stars and found that the observations are consistent with two outflow cavities that are perpendicular to each other. Although the polarization data of IRAS 04016+2610 is still in preparation (Cotera et al. 2006), preliminary analysis of that data suggests that the system may contain two thick, dark lanes which are nearly perpendicular. If our preliminary results are verified, they will provide a confirmation of the models of Wood et al. (2001).

Between the first and second epochs, we find that the overall total observed flux based on the fluxes presented in Table 1, decreases by 44%. Since the 1997 observations do not extend out as far as the 2004 data, if we take irregular but identical apertures which include all of the flux from the 1997 observations, we find that the overall flux decrease is nearly 50%. This is a significantly larger –15– variation than the lower resolution NIR observations by Park & Kenyon (2002), which showed a variability of ∼20%. If we look only at the core of the point source (within the first Airy ring of the PSF), we find the emission decreases by 38%.

Normalizing the total flux to the first epoch (see differenced image in Fig. 11), enables us to investigate any morphological changes in the relative brightness distribution. In the differenced image, we can clearly see a “V” shape in the scattered light, with the northern side decreasing in brightness while the southern side brightens. We can also see that in 2004, there was a bright blob just to the northeast of the star, similar (although much less pronounced) to that seen in DG Tau B (see below), which was not there in 1997. When the images are blinked, we observed changes in the scattered light pattern which appear to be related to differences seen in the immediate vicinity of the star; however, we are unable to discriminate between variability which may caused by a close, ∼20 AU binary and that resulting from a non-uniform stellar photosphere.

3.5. Haro 6-5B

Haro 6-5B (FS Tau B, HBC 381), is located ∼20 west of the close binary FS Tau A. H6-5B is the driving source for the jet HH 157 (Mundt et al. 1984), which extends ∼6 to the northeast, and ∼0.5 to the southwest. Mundt et al. (1984) derived a PA of the jet of 54o, and found that the northeastern jet is blueshifted, while the shorter, southwest jet is redshifted. HST WFPC2 observations of the FS Tau region revealed a star+disk morphology which is similar to HH 30 (Krist et al. 1998), with the brighter portion of the bipolar reﬂection nebula associated with the direction of the blueshifted jet emission.

Yokogawa et al. (2001) obtained 2 mm observations of H6-5B with the Nobeyma Millimeter Array, which provide a good probe of the dusty circumstellar material. They find a deconvolved 2mm disk with a FWHM of 2.30×0.89 (∼ 320×125 AU, assuming a distance of 140 pc Elias 1978) at a PA= 138o ± 7o. Accounting for observational limitations, Yokogawa et al. (2001) derive a physical disk radius of 309±18 AU with an inclination angle of 67o ± 5o. This is in good agreement with previous HST NICMOS observations (Padgett et al. 1999) which found that the NIR dust lane indicates a disk with a diameter of 600 AU, and a PA of 147o. An inclination angle of 67o ± 5o is also in good agreement with previous observations (Eislöffel & Mundt 1998; Krist et al. 1998). The disk mass derived from the 2 mm data is 0.021 M, which is larger than the 0.01 M derived from the HST WFPC2 observations. Krist et al. (1998) note, however, that the WFPC2 observations are not sensitive to heavier particles that may have settled to the disk midplane, and they rely on an ISM grain distribution to derive the disk mass. Krist et al. (1998) suggest an overall geometry for H6-5B, which invokes the presence of a dark cloud between an observed clumpy reflection nebula R1 (nomenclature from Eislöffel & Mundt 1998, , see Fig. 1) located northeast of the bright bipolar nebula, and the extended northeast jet. Their proposed geometry is invoked to explain the apparent disappearance of the jet just before R1, and its reappearance again down stream, 19 away from H6-5B. –16–

Comparison of the two epochs of HST NICMOS data clearly shows the same type of variable lateral asymmetry seen in HH 30. Overall, there is a decrease in the measured ﬂux of ∼25%. The brighter lobe (A in Fig. 1) of the bipolar nebula is much brighter on the right side during the ﬁrst epoch and on the left side at the second epoch. The right side of lobe A decreases by a factor of ∼10% while the left side increases by ∼22% between the two epochs. There is no similar asymmetric change in the emission in the fainter lobe B, with both sides decreasing by ∼36%.

Interestingly, the R1 nebula shows lateral asymmetry in agreement with the lobe A, with emission decreasing on the right side, while increasing on the left side. In fact, on the left side in the ﬁrst epoch, there is a notable gap between the upper nebula and the R1 region, where the emission becomes very faint, essentially negligible. In the second epoch, however, the emission on the left is nearly contiguous. Although the overall brightness increases and decreases within R1, the observed clumps do not move. The overall changes are consistent with R1 being illuminated by the central source in H6-5B, and being the associated with a stationary dark cloud as suggested by Krist et al. (1998). Since H6-5B is nearly edge-on, the changes in the R1 nebula provide information on the high latitude regions of the stellar photosphere, and therefore favor a high latitude hot spot model.

3.6. DG Tau B

DG Tau B is a very low-luminosity YSO (L∗ =0.8L, Jones & Cohen 1986) located 55 southwest of the more luminous and well studied DG Tau (L∗ =8.0 L, e.g. Cohen & Kuhi 1979; Mundt & Fried 1983; Bacciotti et al. 2000; McGroarty & Ray 2004). DG Tau and DG Tau B were some of the first T-Tauri stars to be associated with optical jets (Mundt & Fried 1983). Recently, McGroarty & Ray (2004) have suggested that DG Tau B is the driving source for two newly discovered HH objects in the region, HH 836 and HH 837; with HH 837 located 10.4fromDG Tau B, approximately 0.7 pc at the distance of the Taurus-Auriga cloud. HST WFPC2 observations of DG Tau B (Stapelfeldt et al. 1997) revealed a compact bipolar nebula, with no directly visible star. The 3.5 cm VLA source (Rodríguez, Anglada, & Raga 1995) is located within the dark lane separating the bipolar nebula seen in the WFPC2 images. Padgett et al. (1999) observed DG Tau B with HST NICMOS, and found the the eastern lobe had a strong “V” shape. The peak of the V is located on the axis of symmetry from the blue-shifted jet (Eislöffel & Mundt 1998), which led Padgett et al. (1999) to conclude that they were observing light scattered off the walls of the blue shifted cavity. The western lobe of the cavity was found to be significantly fainter, with an narrower opening angle; the jet was also seen in the lower nebula (location shown in Fig. 1].

As can be seen in Fig. 13, although the general morphology has remained the same, there are several important differences between the two epochs. As can be seen in Table 1, the eastern lobe (A in Fig. 1) was approximately 2.8 times brighter during the first epoch when compared to the second epoch. This difference is almost entirely due to the change in the northern portion of lobe A. During the first epoch, the northern part of lobe A was nearly twice as bright as the –17– southern portion, while in the second epoch they are approximately equal. In fact, the majority of the differences observed between the two epochs comes from this observed asymmetric brightening. In contrast, the western lobe (B) was only 1.3 times brighter in the first epoch. The observed differences in the fluxes do not appear to be entirely correlated with a change in the stellar flux, since aperture photometry inside the first Airy ring of the point source seen in the image indicates a 40% decrease in the flux between the first and second epochs, as compared to an overall decrease of ∼60%.

There are significant differences in the emission from the jet; however, due to the low S/N in the western lobe in the 2004 observations, tracing the observed knots in the jets is not possible. The knots which were furthest from the source, as observed in (Padgett et al. 1999), are too faint to be detected in our data, however, the individual knots seen in our data are not separable from the general scattered light emission in the 1997 data. We do see motion consistent with well collimated jet outflow, but following specific knots is not possible.

3.7. CoKu Tau/1

CoKu Tau/1 is given a Class II classification by White & Hillenbrand (2004) based on NIR spectroscopy. CoKu Tau/1 is associated with the faint bipolar jet HH 156, which, like DG Tau B, is brighter in the redshifted part of the flow than the blueshifted. The outflow velocity for the jet is also notably low, with a value of +30 km s−1 (e.g. Eislöffel & Mundt 1998). Strom & Strom (1994) derived a spectral type of roughly M2 for the star. Padgett et al. (1999), identified a faint binary companion to CoKu Tau/1, which is also seen in our images (see Fig. 14). In the 1997 observations, the binary separation was determined to be 0.240±0.002, with a PA=110o ± 1o. 4. Using Keck HIRES spectra White & Hillenbrand (2004) give a spectral type of K7+M for the binary pair. Eislöffel & Mundt (1998) estimate a lower limit on the inclination of the circumstellar disk of about 81o. The dark lane in CoKu Tau/1 is the most narrow and faint of all those in the current survey; when combined with the small, low velocity jet, this suggests that the object is the most evolved of those presented here. If the inclination angle estimate is correct, the fact that at 1.6 µm we can see both the primary and secondary star in system with a nearly edge-on disk, suggests that we are seeing a more evolved binary pair which still retains only a small circumstellar disk. This conclusion is also supported by the fact that Osterloh & Bechwith (1995) did not detect CoKu Tau/1 at 1.3mm (Osterloh & Bechwith 1995).

Between the two observed epochs, CoKu Tau/1 had one of the smallest changes in overall emission, increasing by only 20%. As can be seen in the normalized difference image, a bright knot of emission on the left side of the upper reflection nebula (A in Fig. 1) has shifted between the two epochs, moving northward. This is the only significant difference between the epochs, other than the increase in the total emission. Upon closer examination of the binary pair first identified in Padgett et al. (1999), we find that the centroid of the fainter companion (CoKu Tau/1 B), has moved slightly (∼ 4 AU) northwest relative to the centroid of brighter primary star (CoKu Tau/1 –18–

A); the detected spatial shift is at approximately 5σ. For CoKu Tau/1, coincidentally, the exact same HST guide stars were used during both epochs, and the difference in the roll angles between the two observations was 0.o57, therefore the expected relative positional uncertainties between the binary companions is very small. By using centroiding in IDP3 on the two binary companions, we estimate that in 2004, the secondary star is located 0.206±0.008 from the primary star, at a PA=107o ±2o. 1. For comparison, our method yields a binary separation of 0.236 and a PA=111o in the first epoch, in good agreement with Padgett et al. (1999). The total projected movement is ∼0. 03, or around 4 AU at the distance of Taurus. Given the proposed inclination of the circumstellar disk of 81o (Eislöffel & Mundt 1998), and assuming the binary orbit is in the same plane as the disk, the actual motion would be on the order of 25 AU.

The observed differences in the scattered light pattern are entirely consistent with the observed movement of the secondary star: i.e. the scattered light pattern has shifted slightly to the northwest on the same side as the fainter companion. Interestingly, the loop of emission which Padgett et al. (1999) suggested may be associated with the fainter star, has not changed between the two epochs. In the normalized difference image shown in Fig. 14, the loop appears to be brighter in the first epoch, but this is an artifact of the normalization, and the flux levels of the loop between the two epochs is unchanged.

4. Discussion

Of the seven objects which we present here, three are either known or likely binary systems: L1551 IRS 5, IRAS 04016+2610 and CoKu Tau/1. Since the data were taken as part of a larger study of dust grain evolution, we purposefully included binary stars to investigate the effect of binaries on the dust population. For the present study, however, the complexities introduced by the multiple radiation sources means that although the observed variations can inform on the effects of the binaries on the scattered light properties, they are not particularly helpful in discriminating between possible intrinsic variation mechanisms. In addition, as discussed above, the variability of HH 30 has been the subject of several papers on the high resolution scattered light data (Wood & Whitney 1998; Stapelfeldt et al. 1999; Watson et al. 2006), and the present observations cannot add substantially to those studies. Therefore, in the following, we will confine ourselves to discussing the observed variations in IRAS 04302+2247 , Haro 6-5B, and DG Tau B.

We wish to address whether the observed variations are the result of cool spots, hot spots, or variations in extinction close to the star. Cool spots are thought to be produced by stellar magnetic activity analogous to the production of Sun spots (e.g. Bertout 1989; Herbst et al. 1994). Some observational hallmarks of cool spots are amplitude variations of less than a few tenths of a magnitude and periodic light curves which persist for long time scales. Generally, cool spots are associated with WTTS/Class III YSOs, with few CTTS/Class II objects having similar observational signatures. For our objects, a detailed literature search did not uncover any observational campaigns to investigate periodicity in their magnitude variations. We searched all the catalogs –19– available as part of the VizieR service provided by the CDS4 for previous photometry of our objects at B-K wavelengths, and determined that at one or more wavelengths all of our objects do have significant photometric variation, typically on the order of several tenths of magnitude (see also Table 1). Although these previously published magnitudes indicate large variability, they are insufficient to extract any periodicity information. The combination of significant magnitude variations and the fact that all of our objects are Class I-Class II, strongly suggests that the observed variations are not produced by rotating cool spots; although we cannot definitely rule out cool spots without additional observations.

Utilizing synoptic photometric and spectroscopic observations of the Class II object AA Tau, Bouvier et al. (2003) have suggested that inclination of the magnetosphere structure has lead to the development of a non-symmetric warp at the inner edge, which accounts for the observed 8.2 day variability, as opposed to a hotspot model. O’Sullivan, et al. (2005) modeled the variability of AA Tau and also found that a warped disk best reproduces the available data. Watson et al. (2006) have observed HH 30 at additional epochs with HST, and find that both a non-circular hot spot and a warped inner disk are consistent with data. For the objects presented here, Haro 6-5B is the only one that shows variations directly analogous to HH 30. The variations seen in H6-5B are well represented by the models of Wood & Whitney (1998), with the asymmetry in the brighter, northwest, lobe shifting to the southwest, while moving in the opposite direction in the other lobe. In addition, however, we see that the bright reflection nebulae to the northwest of H6-5B (R1 from Eislöffel & Mundt 1998), also shows the same asymmetry as the primary nebular lobe, with a shift in peak of the emission to the southwest. Although a warped inner disk might be able to reproduce the observed changes from scattering off a circumstellar disk, how a warp in the disk could simultaneously alter the scattered light pattern of the R1 nebula is less obvious. As noted above, the variation in the R1 nebula strongly suggests the source of the variation is at high latitudes, which is most consistent with a hot spot model.

Recently, Romanova et al. (2004) performed 3-D MHD simulations of disk accretion to an inclined dipole moment and found that the hot spots produced are not necessarily circularly symmetric, but rather may be bow shaped around the magnetic axis. Their proposed geometries would decrease the searchlight effect in the scattered light pattern proposed in the more simple, single hot spot model of Wood & Whitney (1998). In DG Tau B, however, the model proposed by Wood & Whitney (1998), is consistent with the observations, with the caveat that the intensity of a possible hot spot has decreased significantly between the 1997 and 2004 observations. The disappearance of the strongly beamed light seen in 1997 is less consistent with a cool spot as the possible underlying mechanism, in part because one of the hallmarks of cool spots in CCTs is stability over relatively longer time scales (Herbst et al. 1994). Our data for DG Tau B does not have sufficient S/N in the weaker emission lobe (B in Table 1) to definitively rule out contrasting asymmetry in the bipolar lobes; however, to within the uncertainties no asymmetry is observed in the lobe B. This argues

4Centre de Donnes astronomiques de Strasbourg, http://cdsweb.u-strasbg.fr/ –20– for a single hot spot, rather than the bipolar spots proposed by Romanova et al. (2004), but given the uncertainties is speculative.

In IRAS 04302+2247 , we can deﬁnitely see oppositional asymmetries in the two lobes. The quadrapole symmetry, however, is rather complex, lacking the smooth, decidedly conical shape of the other objects. Wolf et al. (2003) modeled IRAS 04302+2247 and although they successfully match the overall symmetry, the models assume a smooth underlying disk, such that the lumpy structure seen in the images (Fig. 10) is not reproduced. The variable asymmetry is seen in some of the observed clumps, but are not mapped consistently as might be expected. In particular, the B2 lobe (see §3.3 above for nomenclature) which we would predict to be dimmer from a rotating hot spot model, is in fact both dimmer and brighter in part. Given the likelihood of extremely patchy extinction, as suggested by the clumpy morphology, with only the variability observed in two epochs, we cannot discount that the variations are the result of equally irregular changes in the extinction near the star. Therefore, the observed variability in IRAS 04302+2247 could be produced by either a warped disk or a rotating hot spot; from the limited data we have available both are plausible.

5. Conclusions

We have presented data on two epochs of HST/NICMOS observations of seven Class I-II YSOs. All of the objects show extensive variability in their scattered light patterns. For three of the objects, the primary cause of the variation is most likely changes in relative positions of the binary stars. For HH 30, we see variability consistent with previous results. For the remaining objects, intermittent hot spots, resulting from magnetic accretion onto the photosphere, provide a plausible explanation for the observed variations, although we cannot preclude variations in the extinction such as a warped inner disk, as a variability mechanism. Additional multi-epoch studies and detailed models are needed to test the proposed variability mechanisms. We will present models of these objects in future papers which will include models of the polarimetry data also obtained as part of this project.

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This preprint was prepared with the AAS LATEX macros v5.2. –25–

Table 1. Data Summary

Observation Log Flux Bipolar Nebulaa (mJy)

Object PID Date Time mJy MF 160W AB

L1551 IRS5 7228 16Nov1997 64s 43.3 11.00 ··· ··· 10178 07Oct2004 192s 31.5 11.34 ··· ··· IRAS 04016+2610 7418 20Dec1997 768s 14.7 12.17 13.2 1.5 10178 24Aug2004 192s 8.2 12.80 7.2 0.8 IRAS 04302+2247 7418 19Aug1997 512s 14.3 12.20 7.7 6.8 10178 24Aug2004 192s 10.5 12.53 4.9 5.8 HH 30 7228 29Sep1997 252s 1.86 14.41 1.40 0.38 10178 07Oct2004 192s 1.92 14.38 1.52 0.28 Haro 6-5B 7418 29Oct1997 512s 10.6 12.52 7.9 0.8 10178 17Dec2004 192s 7.9 12.84 6.3 0.4 Coku Tau/1b 7418 19Aug1997 512s 27.7 11.48 8.8 4.3 10178 25Aug2004 192s 33.2 11.28 11.2 4.3 DG Tau B 7418 21Dec1997 512s 4.9 13.36 3.5 0.63 10178 25Aug2004 192s 2.0 14.33 1.3 0.48 a See Fig. 1 for nomenclature.

b Calculated nebular emission excludes point source to the ﬁrst Airy ring. –26–

Table 2. Astrometry for L1551 IRS 5

Coordinates (J2000) Adjusted Coordinatesa Observations RA Dec RA Dec Authorsb Date λ 4h 31m σ 18o 8 σ 4h 31m σ 18o 8 σ

BH81 23Oct79 ∼10 µm 34.10 0.14 8.03 2.0 34.11 0.14 7.70 2.0 C82 22Mar81 6 cm 34.126 0.006 5.03 0.04 34.137 0.011 4.73 0.15 CS83 5Dec81 2 µm 34.13 0.02 3.99 0.5 34.13 0.02 3.99 0.66 BC85 23Nov83 2 cm 34.135 0.002 5.39 0.03 34.144 0.008 5.14 0.12 34.131 0.002 5.11 0.03 34.140 0.008 4.86 0.12 R86 5Jan85 2 cm 34.137 0.002 5.20 0.03 34.145 0.007 4.97 0.11 34.135 0.002 4.92 0.03 34.143 0.007 4.69 0.11 L95 15Oct91 780 µm 34.10 ··· 4.80 ··· 34.11 ··· 4.69 ··· L97 1Mar96 2.7 mm 34.143 0.009 5.09 0.14 34.144 0.009 5.06 0.14 34.142 0.009 4.72 0.14 34.143 0.009 4.69 0.14 R98 10Jan97 7 mm 34.142 0.002 5.12 0.03 34.143 0.002 5.10 0.03 34.142 0.002 4.80 0.03 34.143 0.002 4.79 0.03 1997 16Nov97 34.141 0.006 5.07 0.08

aWith proper motions, adjusted to 16Nov97.

bBH81: Biechman & Harris 1981. C82: Cohen et al. 1982. CS83: Cohen & Schwartz 1993. BC85: Bieging & Cohen 1985. R86: Rodriquez et al. 1986. L95: Ladd et al. 1995. L97: Looney et al. 1997. R98: Rodriquez et al. 1998.

Table 3. IRAS 04302+2247

F160W Flux (mJy) Source A1 A2 B1 B2

1997 2.44 4.41 3.12 2.97 2004 1.58 2.95 3.21 2.25 2004/1997 0.65 0.67 1.03 0.76 Lucas & Roche (1997) P4 P3 P1 P2 –27–

Fig. 1.— Contour plots of the F160W images for the objects presented in the paper. All of the data presented here are in celestial coordinates The nomenclature used in Table 1 and throughout –28–

Fig. 2.— HH 30. Upper left: Square root scaled F160W image from the 1997 observations. Upper right: Square root scaled F160W of the 2004 observations. Lower The diﬀerenced image, 2004-1997. –29–

Fig. 3.— Three color composite of L1551 IRS5. F110W is blue, F160W is green, and F204M is red. The northern jet and the region “D” are seen as blue-green features since they are not apparent in the F204M image. See Fig. 4 for nomenclature. . –30–

Fig. 4.— L1551 IRS 5 F160W contour map (1997). The contour levels were selected to illustrate both the faint and strong features clearly, and do not increase at uniform intervals. The characteristic parabolic pattern of light scattering off a flared optically thick circumstellar disk is indicated, as is the location of the “wake” surrounding a clump of emission to the north northwest of the main portion of the reflection nebula. –31–

Fig. 5.— Linecuts through the peak of the compact source along a north northeast to south southwest axis, along the ridge of the reﬂected emission. The dashed lines are TinyTIM PSF models for these observations at the various wavelengths. –32–

Fig. 6.— Left: Astrometrically correct F160W contours overlaid with published position of the binary sources in L1551 IRS 5. The plus signs indicate positions from BC85, astericks are from R86, diamonds are from L97, and triangles are from R98 (see Table 2). Right: The same observations, precessed to the date of the NICMOS observations, using the proper motion value for IRS 5 derived here. The box gives the 2σ position of the compact NICMOS source. As can be seen, the position of the IRS 5A source is well correlated with the NICMOS source. –33–

Fig. 7.— L1551 IRS 5 extinction map. Solid lines are in AV , dashed lines are the flux contours from the F160W filter image, X marks the location of the NIR peak and clump, while the boxes indicate the 1σ positions of IRS 5A and 5B based on the positions of Rodríguez et al. (1998). The extinction is based on the F160W-F204M colormap, and the extinction law of Rieke & Lebofsky (1985). –34–

Fig. 8.— L1551 IRS 5. Upper left: Log scale F160W image from the cycle 7 observations. The “C”pattern, indicative of scattering off a circumstellar disk, can be seen, as can the outflow from the jet. Upper right: Log scale F160W cycle 14 observations. In addition to the features seen in the early epoch, the jet from the binary companion is faintly visible. Lower: The differenced image, 2004-1997. The jet seen in 1997 has clearly propagated outward between the two epochs. The differences between the two epochs are consistent with orbital motion of a binary star. –35–

Fig. 9.— Upper: Contour plots of the jet, with the levels selected to highlight the features of interest. The plots have been overlaid with a trace of the peak emission along the jet. Lower: Grayscale images of the same regions as above. Here we have placed the traces for both jets on each image (green 1997, red 2004). On the left we illustrate the propagation of the jet outward from the main source. On the right we have aligned the inﬂection knots (by a simple translation of the jet trace), to illustrate the rotation of the jet emission relative to inﬂection point. In both the upper and lower panels, the 1997 data is on the left, and the 2004 data is on the right. –36–

Fig. 10.— IRAS 04302+2247. Upper left: Log scale F160W image from the cycle 7 observations. Upper right: Log scale F160W cycle 14 observations. Lower: The diﬀerenced image, 2004-1997. –37–

Fig. 11.— IRAS 04016+2610. Upper left: Log scale F160W image from the cycle 7 observations. Upper right: Log scale F160W cycle 14 observations. Lower: The diﬀerenced image, 2004-1997. –38–

Fig. 12.— Haro 6-5B (FS Tau B). Upper left: Log scale F160W image from the cycle 7 observations. Upper right: Log scale F160W cycle 14 observations. Lower: The diﬀerenced image, 2004-1997. –39–

Fig. 13.— DG Tau B. Upper left: Log scale F160W image from the cycle 7 observations. Upper right: Log scale F160W cycle 14 observations. Lower: The diﬀerenced image, 2004-1997. –40–

Fig. 14.— CoKu Tau/1. Upper left: Log scale F160W image from the cycle 7 observations. Upper right: Log scale F160W cycle 14 observations. Lower: The diﬀerenced image, 2004-1997.