\ion{He}{1} 1.083 $\mu$m emission and absorption in DG Tau: line excitation in jet, hot wind, and accretion flow\footnote{Based on data collected at Subaru Telescope, which is operated by the National Astronomical Observatory of Japan.}
M. Takami\altaffilmark{2}
A. Chrysostomou\altaffilmark{2}
J. Bailey\altaffilmark{3}
T.M. Gledhill\altaffilmark{2}
M. Tamura\altaffilmark{4}
and H. Terada\altaffilmark{5}
We present long-slit spectroscopy and spectro-astrometry of \ion{He}{1} 1.083 $\mu$m emission in the T Tauri star, DG Tau. We identify three components in the \ion{He}{1} feature: (1) a blueshifted emission component at $v$$\simeq$$-$200 km s$^{-1}$, (2) a bright emission component at zero-velocity with a FWZI of $\sim$500 km s$^{-1}$, and (3) a blueshifted absorption feature at velocities between $-$250 and $-$500 km s$^{-1}$. The position and velocity of the blueshifted \ion{He}{1} emission coincide with a high-velocity component (HVC) of the [\ion{Fe}{2}] 1.257 $\mu$m emission, which arises from a jet within an arcsecond of the star. The presence of such a high excitation line (excitation energy $\sim$ 20 eV) within the jet supports the scenario of shock heating. The bright \ion{He}{1} component does not show any spatial extension, and it is likely to arise from magnetospheric accretion columns. %The blueshifted absorption does not show the presence of the associated off-stellar emission, suggesting that the \ion{He}{1} wind in DG Tau has a spatial scale well below our spatial resolution. The blueshifted absorption shows greater velocities than that in H$\alpha$, suggesting that these absorption features arise from the accelerating wind close to the star.
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\shorttitle{\ion{He}{1} emission and absorption in DG Tau}
\shortauthors{Takami et al.}

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\title{\ion{He}{1} 1.083 $\mu$m emission and absorption in DG Tau: line excitation in jet, hot wind, and accretion flow\footnote{Based on data collected at Subaru Telescope, which is operated by the National Astronomical Observatory of Japan.}}

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\author{M. Takami\altaffilmark{2}, A. Chrysostomou\altaffilmark{2}, J. Bailey\altaffilmark{3}, T.M. Gledhill\altaffilmark{2}, M. Tamura\altaffilmark{4}, and H. Terada\altaffilmark{5}}

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%\altaffiltext{1}{Based on data collected at Subaru Telescope, which is operated by the National Astronomical Observatory of Japan.} 
\altaffiltext{2}{Department of Physical Sciences, University of Hertfordshire, College Lane, Hatfield Herts, AL10 9AB, UK} 
\altaffiltext{3}{Anglo-Australian Observatory, PO Box 296, Epping, NSW 1710, Australia} 
\altaffiltext{4}{National Astronomical Observatory of Japan, Osawa, Mitaka, Tokyo 181-8588, Japan} 
\altaffiltext{5}{Subaru Telescope, National Astronomical Observatory of Japan, 650 North A'ohoku Place, Hilo, Hawaii 96720} 
\email{takami@star.herts.ac.uk}

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\begin{abstract}
We present long-slit spectroscopy and spectro-astrometry of \ion{He}{1} 1.083 $\mu$m emission in the T Tauri star, DG Tau. We identify three components in the \ion{He}{1} feature: (1) a blueshifted emission component at $v$$\simeq$$-$200 km s$^{-1}$, (2) a bright emission component at zero-velocity with a FWZI of $\sim$500 km s$^{-1}$, and (3) a blueshifted absorption feature at velocities between $-$250 and $-$500 km s$^{-1}$. The position and velocity of the blueshifted \ion{He}{1} emission coincide with a high-velocity component (HVC) of the [\ion{Fe}{2}] 1.257 $\mu$m emission, which arises from a jet within an arcsecond of the star.
The presence of such a high excitation line (excitation energy $\sim$ 20 eV) within the jet supports the scenario of shock heating. The bright \ion{He}{1} component does not show any spatial extension, and it is likely to arise from magnetospheric accretion columns.
%The blueshifted absorption does not show the presence of the associated off-stellar emission, suggesting that the \ion{He}{1} wind in DG Tau has a spatial scale well below our spatial resolution.
 The blueshifted absorption shows greater velocities than that in H$\alpha$, suggesting that these absorption features arise from the accelerating wind close to the star.
\end{abstract}

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\keywords{accretion, accretion disks---line: formation---stars: individual (DG Tau)---stars: pre-main sequence---ISM: jets and outflows}

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\section{Introduction}
Understanding the mechanisms of mass accretion and the driving of jets/winds is one of the most important key issues of star formation theories. Disk accretion rates estimated through UV excess measurements and/or permitted line luminosities are correlated with outflow signatures such as forbidden line luminosities (e.g., Hartigan, Edwards, \& Ghandour 1995), suggesting a physical link between the two phenomena: i.e., jets/winds remove excess angular momentum from accreting material, thereby preventing young stellar objects from spinning up to breakup velocity.
Several theories have been proposed to attempt to explain this physical link (e.g., Shu et al. 2000; K\"onigl \& Pudritz 2000; Goodson, B\"ohm \& Winglee 1999), however, consensus has not yet been achieved. 

Low-mass pre-main sequence stars (T Tauri stars) exhibit a rich variety of emission lines at optical-to-IR wavelengths, probing the activities of jets, winds, and accretion flow close to the stars. 
%The densities and temperatures required to excite permitted lines indicate formation within $\sim$10 R$_*$ of the stellar surface, 
Permitted atomic emission lines have often been considered to arise from magnetospheric accretion columns and/or accretion shocks. Indeed, these line luminosities show correlation with the accretion luminosity (e.g., Muzerolle, Hartmann, \& Carvet 1998a), and idealized magnetospheric accretion models have been successful in reproducing at least some of the observed permitted line profiles (Hartmann, Hewett, \& Carvet 1994; Muzerolle, Calvet, \& Hartmann 1998b, 2001). On the other hand,
%a significant contribution from jets/winds has been suggested in some T Tauri stars via the measurements of line profiles and/or spatial extension at AU scales (e.g., Beristain et al. 2001; Takami et al. 2001).
H$\alpha$ and \ion{He}{1} profiles in some T Tauri stars exhibit blue-shifted wings, and Beristain, Edwards, \& Kwan (2001) argue that a hot wind can also contribute to their emission. Spectro-astrometric observations by Takami et al. (2001) have revealed that a bipolar outflow at AU scales is responsible for the H$\alpha$ emission in the T Tauri star RU Lupi.
Forbidden lines in T Tauri stars arise from outflowing gas at scales of more than 1 AU, and often exhibit two blueshifted components at high and low velocities (e.g., Hirth, Mundt, \& Solf 1997). Different densities and temperatures for these two components is suggestive of their distinct origins, e.g., a fast jet and a slow disk wind (see Eisl\"offel et al. 2000).
%The LVC velocity in T Tauri stars increased for line with decreasing critical density, suggesting that expected for accelerating wind (Hirth, Mundt, \& Solf 1994; Hartigan et al. 1995).

We present results of long-slit spectroscopy and spectro-astrometry for \ion{He}{1} 1.083 $\mu$m and [\ion{Fe}{2}] 1.257 $\mu$m emission in DG Tau. The former is one of the brightest among \ion{He}{1} lines at optical-to-IR wavelenths, and its diagnostic potential lies in the high-excitation energy (20 eV), which restricts its formation to a region either of high temperature or close proximity to a source of ionizing radiation (cf. Beristain et al. 2001). DG Tau is one of the most active young stellar objects known.  Analysis of the bright, hot continuum excess gives a mass accretion rate of 5$\times$10$^{-7}$ M$_{\odot}$ yr$^{-1}$ (Gullbring et al. 2000), one of the largest among T Tauri stars in the Taurus-Auriga cloud (Muzerolle et al. 1998a). Beristain et al. (2001) show a blue-shifted wing to the \ion{He}{1} 5876 {\AA} emission at velocities up to 600 km s$^{-1}$, indicating the presence of a hot wind close to the star. Recent high resolution observations have revealed knotty, bubble-like, and bow-shape structures in the associated jet within a few arcsecond of the object (e.g., Kepner et al. 1993; Dougados et al. 2000; Bacciotti et al. 2000). In this Letter, we show that the \ion{He}{1} 1.083 $\micron$ feature in DG Tau consists of the following three components: a blueshifted {\it emission} component from a jet, a blueshifted {\it absorption} component from a hot wind, and a bright emission component at zero-velocity presumably from magnetospheric accretion columns.

%\newpage
\section{Observations}
Observations were carried out on 2000 December 13 at the SUBARU 8.2-m
telescope using the Infrared Camera and Spectrograph IRCS (Tokunaga et al. 1998; Kobayashi et al. 2000). The echelle grating mode with a 0.3-arcsec wide slit provides a spectral resolution
of 10$^4$. The pixel scale of 0.075 arcsec provides good sampling of the seeing profile (FWHM$\sim$1 arcsec during our observations), thereby allowing accurate spectro-astrometric measurements (Bailey 1998a,b). The spectra were obtained at four slit position angles (0, 90, 180, and 270$^\circ$). The DG Tau jet lies at a position angle of $\sim 225^\circ$ (Lavalley et al. 1997; Eisl\"{o}ffel \& Mundt 1998), thus the selected slit angles minimize contamination from the extended jet and allows us to study the emission line regions closest to the star. In addition to the target, an A-type bright standard was observed at similar airmasses to correct for telluric absorption. The flat fields were made by combining many exposures of the spectrograph illuminated by a halogen lamp. 

The data were reduced using the KAPPA and FIGARO packages provided by Starlink. The position-velocity diagrams were obtained via standard reduction processes: dark-subtraction, flat-fielding, removal of bad pixels, correcting for curvature in the echelle spectra, wavelength calibration, correcting for telluric absorption, and night-sky subtraction. Wavelength calibrations were made using atmospheric absorption features on the stellar continuum via comparison with modelled spectra provided by ATRAN (Lord 1992). The systemic motion of the object was calibrated 
based on previous observations of \ion{Li}{1} 6707 {\AA} photospheric absorption (V$_{\rm{Hel}}$=+16.5 km s$^{-1}$--- Bacciotti et al. 2000).

In addition to the position-velocity diagrams, spectro-astrometric ``position spectra'' were obtained to study the spatial structure of emission line regions down to milliarcsecond scales (Bailey 1998a,b; Takami et al. 2001). These spectra were determined after dark-subtraction and flat-fielding by fitting the seeing profile at each wavelength with a Gaussian function.
About 6$\times$10$^4$ photons at each wavelength provide a typical accuracy of 6 milliarcsec at the continuum level.
Any instrumental effects in the position spectra were eliminated with high accuracy by subtracting those with opposite position angles ($0 - 180 ^\circ$ or $90 - 270 ^\circ$).
%, since they will remain constant on rotation whereas any real structure in the source will reverse its sign. %Effects which are removed in this way include any curvature or optical distortion introduced by the spectrograph, any misalignment of the spectrum with the detector columns, any effects due to departure of the detector pixels from a regular grid, and imperfect flat fielding.
Each position spectrum has an arbitrary zero point, which is adjusted to correspond to the continuum position. To determine the true-position of the line emission, contamination by the continuum was removed using the line-to-continuum ratio at each wavelength (see Takami et al. 2001 for detail).
%the observed position was corrected for the effect of the underlying continuum by dividing the positional displacement in the raw spectra by $I_{\lambda \rm{(line)}}/(I_{\lambda \rm{(line)}}+I_{\lambda \rm{(conn)}})$.

\section{Results}
Fig. 1 shows the position-velocity diagrams for \ion{He}{1} 1.083 $\mu$m and [\ion{Fe}{2}] 1.257 $\mu$m emission along north-south (N-S) and east-west (E-W) slit positions.
The figure exhibits two components in both the \ion{He}{1} and [\ion{Fe}{2}] emission.
The \ion{He}{1} emission consists of a bright component centered on the systemic velocity, and a high velocity component (HVC) at a velocity of roughly $-200$ km s$^{-1}$.
The former component does not show any spatial displacement from the star, whose position is indicated by the dot-dashed line in the figure.
On the other hand, the HVC extends towards the south and west peaking 0".4 south of the star.
The position of the HVC coincides with structure in the jet located $\sim$0".6 southwest from the star, previously identified by Bacciotti et al. (2000).

The position-velocity diagrams of the [\ion{Fe}{2}] emission exhibit two blueshifted velocity components as observed in optical forbidden lines of many T Tauri stars (e.g., Hirth et al. 1997).
These components in our results exhibit velocities of $-$200 to $-$250 km s$^{-1}$ and $-$100 to $-$150 km s$^{-1}$, and their centroidal positions are displaced from the star in the direction of the jet by $\sim$0".4 arcsec and $\sim$0".2 arcsec, respectively, in both the N-S and E-W diagrams. %\vspace{0.3cm}
The position and the velocity of the HVC in the [\ion{Fe}{2}] emission coincide with those of the \ion{He}{1} emission described above. 

Fig. 2 shows intensity profiles in the regions (a)(b)(c) given in Fig. 1: each 0", 0".4, and 0".8 arcsecs from the star in the north-south diagrams, respectively. The \ion{He}{1} 1.083 $\mu$m profile in region (a) exhibits a triangular component at zero velocity with a FWZI of $\sim$500 km s$^{-1}$, corresponding to the bright component in Fig. 1.
%The profile also shows marginal excess emission at $v$$\sim$$-$200 km s$^{-1}$, corresponding to the HVC.
The HVC clearly appears in profiles (b) and (c), exhibiting a peak at $-$205 km s$^{-1}$.
In addition to these emission components, profile (a) shows blueshifted absorption between $-$250 and $-$500 km s$^{-1}$.
The intensity profiles in [\ion{Fe}{2}] emission in Fig. 2 exhibit blueshifted peaks at $-$230 to $-$240 and $-$105 km s$^{-1}$, respectively. These [\ion{Fe}{2}] profiles are similar to those of optical emission lines observed by Hamann (1994), especially of [\ion{Fe}{2}] 8617 {\AA}, 7452 {\AA}, and 7388 {\AA} lines.

Fig. 2 also shows position spectra obtained via spectro-astrometric reduction. Displacement is not detected in the brightest emission component of the \ion{He}{1} emission within the measured positional accuracy of $\sim$10 milliarcsec, corresponding to $\sim$1.4 AU at the observed star-forming region (Wichmann et al. 1998). This is consistent with the argument that permitted atomic emission lines in T Tauri stars form within $\sim$10 R$_*$ of the stellar surface (e.g., Najita et al. 2000). The position spectra in the low-velocity component (LVC) of the [\ion{Fe}{2}] emission show a constant displacement of $\sim$0".2 within the error bars. A constant displacement is also observed in the optical forbidden lines of RU Lupi (Takami et al. 2001), and contrasts with that expected in an accelerating flow in which the displacement should increase with the velocity. Such constancy in the position spectra could be explained if the line broadening is due to rotation around the flow axis (Hirth, Mundt, \& Solf 1994), although the error bars in Fig. 2 cannot rule out the contribution of accelerating motion. 
%Indeed, Hirth et al. (1994) and Hartigan et al. (1995) shows that the LVC velocity in T Tauri stars increased for line with decreasing critical density, consistent with that expected for accelerating wind.
The measured position angles in the emission features are $215\pm 6 ^\circ$, $213 \pm 14 ^\circ$, and $219\pm 17 ^\circ$ for the \ion{He}{1} HVC, [\ion{Fe}{2}] HVC, and [\ion{Fe}{2}] LVC, respectively, coincident with the position angle of the extended jet (Lavalley et al. 1997; Eisl\"{o}ffel and Mundt 1998).

\section{Discussion}
It is surprising that the \ion{He}{1} 1.083 $\mu$m emission is detected within the HVC of the forbidden line region. The HVC of forbidden line regions in T Tauri stars have been considered to arise from a jet close to the source (Eisl\"offel et al. 2000), and mainly observed in [\ion{O}{1}] 6300 {\AA}, [\ion{S}{2}] 6716/6731 {\AA}, and [\ion{N}{2}] 6583 {\AA}, whose excitation energies are typically 2 eV from the ground. On the other hand, the \ion{He}{1} 1.083 $\mu$m emission has an excitation energy of 20 eV, about $\sim$10 times larger. This line has been observed in some high-velocity Harbig-Haro objects (e.g., Brugel, B\"ohm, \& Mannery 1981) and in \ion{H}{2} regions (e.g., Takami et al. 2002), however, it is not easily thermally excited, requiring conditions with temperatures of $\sim$10$^4$ K and electron densities of $<$10$^6$ cm$^{-3}$.

The heating mechanism in the jet close to the source has been debated over many years. Proposed heating mechanisms include: shocks caused by variations in the velocity of the jet (e.g., Hartigan et al. 1995), ambipolar diffusion (Safier 1993; Garcia 2001), and turbulent dissipation in a viscous mixing-layer (e.g., Binette et al. 1999). Recent high-resolution imaging and emission line diagnostics have prefered shock heating (Bacciotti et al. 2000; Lavalley-Fouquet, Cabrit, \& Dougados 2000), and our results support this conclusion for the following reasons. One of the outer working surfaces of the DG Tau jet exhibit [\ion{O}{3}] 4959 {\AA} and 5007 {\AA} emission (Cohen \& Fuller 1985), indicating that O$^+$ ions are ionized beyond their ionization potential of 35 eV: thus, shock velocities in the DG Tau jet should be sufficient to excite the \ion{He}{1} 1.083 $\mu$m emission, the excitation energy of which is 20 eV.
%The position and velocity of the observed \ion{He}{1} coincide with those of the [\ion{Fe}{2}] 1.257 $\mu$m emission, agreeing with the prediction of shock models that various ionization conditions can be achieved in the geometrically thin postshock regions (e.g., Raymond 1979).
On the other hand, the models of ambipolar diffusion and turbulent dissipation predict temperatures of up to $\sim$$3 \times 10 ^4$ K (Safier 1993; Garcia et al. 2001; Binette et al. 1999) which, together with the inferred electron densities ($\sim$10$^5$ cm$^{-3}$ --- Lavalley-Fouquet et al. 2000), is not likely to explain the presence of such a highly excited line.

The velocities of the \ion{He}{1} blueshifted $absorption$ roughly coincide with the blueshifted wing $emission$ of \ion{He}{1} 5576 {\AA}, which extends up to 600 km s$^{-1}$ (Beristain et al. 2001). This difference is attributed to their different optical thickness. The lower energy level of the \ion{He}{1} 1.083 $\mu$m transition is ``metastable'', and its large population causes an absorption feature via resonance scattering in the outer region of a non-spherical wind. On the other hand, the lower level of the \ion{He}{1} 5576 {\AA} transition must be much less populated, allowing the emission from the inner wind to pass through the outer region without absorption. 
%Beristain et al. (2001) argue that their \ion{He}{1} 5576 {\AA} profiles exhibit a blue-shifted absorption at approximately $-$200 km s$^{-1}$. However, our 1.083 $\mu$m profile at the star does not show any clear evidence of the absorption at this velocity despite its larger optical thickness. We thus conclude that the dip in their \ion{He}{1} 5576 {\AA} profile is not due to the absorption in a wind but an intensity minima between the bright and blue-shifted emission components.
Our results do not show the presence of off-stellar \ion{He}{1} emission associated with the blueshifted absorption, as observed in a massive star $\theta ^1$C Ori (Oudmaijer et al. 1997). This fact suggest that the \ion{He}{1} wind in DG Tau has a spatial scale well below our spatial resolution. Interestingly, the H$\alpha$ profiles observed by Beristain et al. (2001) exhibit a blueshifted absorption at much lower velocities, peaking at roughly $-$50 km s$^{-1}$. Such different velocities between the \ion{He}{1} 1.083 $\mu$m and H$\alpha$ absorption may be a result of acceleration in the wind. 

The bright \ion{He}{1} emission component at zero velocity appears both in \ion{He}{1} 1.083 $\mu$m and 5876 {\AA}, in contrast with the wind at $-250$ to $-500$ km s$^{-1}$. Such different line excitation between the two components suggest distinct origins. A possible explanation is that the bright component arises in magnetospheric accretion columns, as often argued for permitted lines. Indeed, the profiles of this component are similar to those of Pa$\beta$ (Folha \& Emerson 2001), whose luminosity show a tight correlation with the accretion luminosity (Muzerolle et al. 1998a). 
%In addition, the \ion{He}{1} 1.083$\mu$m profile in Fig. 2 appears to have a red-shifted absorption, and such a triangular profile with a red-shifted absorption is similar to those predicted by magnetospheric accretion models for optically-thick lines (Hartmann et al. 1994). 

%Neither the position-velocity diagrams nor position spectra exhibit any spatial extension in the brightest \ion{He}{1} component. This is consisitent with the \ion{He}{1} line arising from magnetospheric channel flow or a hot wind close to the star, a conclusion also arrive at by Beristain et al. (2001). 
%%The emission component also appears in the \ion{He}{1} 5876 {\AA} and 6678 {\AA} profiles observed by Beristain et al. (2001), and our results are consistent with their conclusion that the \ion{He}{1} lines arise from magnetospheric mass accretion columns and/or a hot wind close to the star.
%%Interestingly, the high velocity wind up to 500-600 km s$^{-1}$ appears in the 1.083 $\mu$m profile as absorption, in their \ion{He}{1} 5876 {\AA} profiles as emission, and is absent in the \ion{He}{1} 6678 {\AA} profile (Beristain et al. 2001). Such different line excitation between the brightest emission component and a high-velocity wind suggests a distinct origin of the former component. Possible explanations include: (1) the component arises from the accretion column, and (2) it arises from the accelerating wind while the blue-shifted absorption/emission arises from the outer region where the acceleration is terminated.\vspace{0.1cm}


\vspace{0.5cm}
The authors thank the anonymous referee for useful comments, and T. Onaka for use of ATRAN on his workstation.
The authors also acknowledge the data analysis facilities provided by the
Starlink Project which is run by CCLRC on behalf of PPARC.  MT thanks PPARC for support through a PDRA.
 
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\begin{thebibliography}{}
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\bibitem[Beristain et al.~2001]{BEK00} Beristain, G., Edwards, S., \& Kwan, J.\ 2001, \apj,  551, 1037 
\bibitem[Binette et al.~1999]{Binette99} Binette, L., Cabrit, S., Raga, A., \& Cant{\' o}, J.\ 1999, \aap,  346, 260
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%\bibitem[Hamann et al.~1994]{HSCP94} Hamann F., Simon M., Carr J.~S., Prato L., 1994, \apj,  436, 292 
\bibitem[Hartigan et al.~1995]{HEG95} Hartigan, P., Edwards, S., \& Ghandour, L.\ 1995, \apj,  452, 736
\bibitem[Hartmann et al.~1994]{H94} Hartmann, L., Hewett, R., \& Calvet, N.\ 1994, \apj,  426, 669
%\bibitem[Hirth, Mundt, \& Solf(1994)]{1994A&A...285..929H} Hirth, G.~A., Mundt, R., \& Solf, J.\ 1994, \aap, 285, 929
\bibitem[Hirth, Mundt, \& Solf(1997)]{1997A&AS..126..437H} Hirth, G.~A., Mundt, R., \& Solf, J.\ 1997, \aaps, 126, 437
\bibitem[Kepner et al.~1993]{Kepner93} Kepner, J., Hartigan, P., Yang, C., \& Strom, S., 1993, \apjl,  415, L119
%\bibitem[Kitamura et al.~1996]{Kiramura96} Kitamura, Y., Kawabe, R., \& Saito, M.\ 1996, \apj,  457, 277
\bibitem[Kobayashi et al.(2000)]{2000SPIE.4008.1056K} Kobayashi, N.~et al.\ 2000, \procspie, 4008, 1056
\bibitem[Konigl \& Pudritz(2000)]{2000prpl.conf..759K} Konigl, A.~\& Pudritz, R.~E.\ 2000, Protostars and Planets IV (Book - Tucson: University of Arizona Press; eds Mannings, V., Boss, A.P., Russell, S.~S.), p.~759, 759
\bibitem[Lavalley et al.(1997)]{1997A&A...327..671L} Lavalley, C., Cabrit, S., Dougados, C., Ferruit, P., \& Bacon, R.\ 1997, \aap, 327, 671
\bibitem[Lavalley-Fouquet, Cabrit, \& Dougados(2000)]{2000A&A...356L..41L} Lavalley-Fouquet, C., Cabrit, S., \& Dougados, C.\ 2000, \aap, 356, L41
\bibitem[Lord(1992)]{lor92} Lord, S. D.\ 1992, NASA Technical Memor. 103957 
\bibitem[Muzerolle et al.~1998]{Muzerolle98a} Muzerolle, J., Calvet, N., \& Hartmann, L.\ 1998b, \apj,  492, 743
\bibitem[Muzerolle et al.~2001]{Muzerolle01} Muzerolle, J., Calvet, N., \& Hartmann, L.\ 2001, \apj,  550, 944
\bibitem[Muzerolle et al.~1998]{Muzerolle98b} Muzerolle, J., Hartmann, L., \& Calvet, N., 1998a, \aj,  116, 2965
\bibitem[Najita et al.~2000]{N2000} Najita, J.~R., Edwards, S., Basri, G., \& Carr J., 2000  Protostars and Planets IV (Book - Tucson: University of Arizona Press; eds Mannings, V., Boss, A.P., Russell, S.~S.), p.~457
\bibitem[Oudmaijer, Busfield, \& Drew(1997)]{1997MNRAS.291..797O} Oudmaijer, R.~D., Busfield, G., \& Drew, J.~E.\ 1997, \mnras, 291, 797
\bibitem[Safier 1993a]{Safier93a} Safier, P.~N.\ 1993, \apj,  408, 115
\bibitem[Shu et al.~2000]{Shu00} Shu, F.~H., Najita, J.~R., Shang, H., \& Li, Z.-Y., 2000\ , Protostars and Planets IV (Book - Tucson: University of Arizona Press; eds Mannings, V., Boss, A.P., Russell, S.~S.), p.~789
\bibitem[Takami et al.(2001)]{2001MNRAS.323..177T} Takami, M., Bailey, J., Gledhill, T.~M., Chrysostomou, A., \& Hough, J.~H.\ 2001, \mnras, 323, 177
\bibitem[Takami et al.~2002]{Takami02} Takami, M. et al. 2002, \apj, 566, 910
\bibitem[Tokunaga et al.(1998)]{1998SPIE.3354..512T} Tokunaga, A.~T.~et al.\ 1998, \procspie, 3354, 512
\bibitem[Wichmann et al.(1998)]{1998MNRAS.301L..39W} Wichmann, R., Bastian, U., Krautter, J., Jankovics, I., \& Rucinski, S.~M.\ 1998, \mnras, 301, L39
\end{thebibliography}


\clearpage

%% Use the figure environment and \plotone or \plottwo to include 
%% figures and captions in your electronic submission.

\begin{figure*}
\plotone{f1.eps}
%\plotone{contour.eps}
\caption{Continuum-subtracted position-velocity diagrams for the \ion{He}{1} 1.083 $\mu$m and [\ion{Fe}{2}] 1.257$\mu$m lines in the north-south and east-west directions. The spacing of the contour corresponds to 5\% and 10\% of the peak flux in the \ion{He}{1} and [\ion{Fe}{2}] diagrams, respectively.
A dot-dashed line in each figure shows the position of the star. Scale bars in the bottom-right figure show the average seeing size and the spectral resolution. Intensity profiles are extracted from the regions (a)(b)(c), and shown in Figure 2.
Excess emission between dotted lines may be due to imperfect sky subtraction. \label{fig1}}
\end{figure*}

\clearpage 

\begin{figure*}
\plotone{f2.eps}
%\plotone{Fig2_revised2.eps}
\caption{Intensity and position spectra of the \ion{He}{1} 1.083 $\mu$m and [\ion{Fe}{2}] 1.257$\mu$m lines. Solid, dotted, and dashed intensity profiles were obtained from the regions (a)(b)(c) in the N-S diagrams (see Fig. 1), and normalized to the continuum flux in (a). Position spectra were obtained from the whole seeing profile including these three regions. The narrow hump in the \ion{He}{1} profiles indicated by the arrow may be due to imperfect sky subtraction. Position spectra are shown in the velocity ranges in which spatial structure is marginally or negatively shown in Fig. 1. \label{fig2}}
\end{figure*}

%% If you are not including electonic art with your submission, you may
%% mark up your captions using the \figcaption command. See the 
%% User Guide for details.
%%
%% No more than seven \figcaption commands are allowed per page, 
%% so if you have more than seven captions, insert a \clearpage 
%% after every seventh one. 

%% Tables should be submitted one per page, so put a \clearpage before
%% each one.

%% Two options are available to the author for producing tables:  the
%% deluxetable environment provided by the AASTeX package or the LaTeX
%% table environment.  Use of deluxetable is preferred.
%%

%% Three table samples follow, two marked up in the deluxetable environment,
%% one marked up as a LaTeX table.

%% In this first example, note that the \tabletypesize{}
%% command has been used to reduce the font size of the table.
%% Note also that the \label command needs to be placed 
%% inside the \tablecaption.

%% Text for table notes should follow after the \enddata but before
%% the \end{deluxetable}. Make sure there is at least one \tablenotemark
%% in the table for each \tablenotetext.

%% If you use the table environment, please indicate horizontal rules using
%% \tableline, not \hline.
%% Do not put multiple tabular environments within a single table.
%% The optional \label should appear inside the \caption command.

%% Any table notes must follow the \end{tabular} command.

%% If the table is more than one page long, the width of the table can vary
%% from page to page when the default \tablewidth is used, as below.  The
%% individual table widths for each page will be written to the log file; a
%% maximum tablewidth for the table can be computed from these values.
%% The \tablewidth argument can then be reset and the file reprocessed, so
%% that the table is of uniform width throughout. Try getting the widths
%% from the log file and changing the \tablewidth parameter to see how
%% adjusting this value affects table formatting.

%% In this example, we have used the optional * argument to \\ to
%% instruct LaTeX to keep rows together on the same page. (See the
%% lines following the \cutinhead.) Using \\* to group together table
%% rows on the same page affects how the table breaks. Try taking
%% the *'s out and LaTeXing again to see the difference.

%% You can append references to a table using the \tablerefs command.

%% Tables may also be prepared as separate files. See the accompanying
%% sample file table.tex for an example of an external table file.
%% To include an external file in your main document, use the \input
%% command. Uncomment the line below to include table.tex in this
%% sample file. (Note that you will need to comment out the \documentclass,
%% \begin{document}, and \end{document} commands from table.tex if you want
%% to include it in this document.)

%% \input{table}
		

%% The following command ends your manuscript. LaTeX will ignore any text
%% that appears after it.

\end{document}

%%
%% End of file `sample.tex'.

astro-ph0202327.extracted_bib
\begin{thebibliography} is followed by an empty set of
%% curly braces.  If you forget this, LaTeX will generate the error
%% "Perhaps a missing \item?".
%%
%% thebibliography produces citations in the text using \bibitem-\cite
%% cross-referencing. Each reference is preceded by a
%% \bibitem command that defines in curly braces the KEY that corresponds
%% to the KEY in the \cite commands (see the first section above).
%% Make sure that you provide a unique KEY for every \bibitem or else the
%% paper will not LaTeX. The square brackets should contain
%% the citation text that LaTeX will insert in
%% place of the \cite commands.

%% We have used macros to produce journal name abbreviations.
%% AASTeX provides a number of these for the more frequently-cited journals.
%% See the Author Guide for a list of them.

%% Note that the style of the \bibitem labels (in []) is slightly
%% different from previous examples.  The natbib system solves a host
%% of citation expression problems, but it is necessary to clearly
%% delimit the year from the author name used in the citation.
%% See the natbib documentation for more details and options.

\begin{thebibliography}{}
\bibitem[Bailey(1998a)]{1998SPIE.3355..932B} Bailey, J.~A.\ 1998, \procspie, 3355, 932
\bibitem[Bailey(1998b)]{1998MNRAS.301..161B} Bailey, J.\ 1998, \mnras, 301, 161
\bibitem[Bacciotti et al.(2000)]{2000ApJ...537L..49B} Bacciotti, F., Mundt, R., Ray, T.~P., Eisl{\" o}ffel, J., Solf, J., \& Camezind, M.\ 2000, \apjl, 537, L49
\bibitem[Beristain et al.~2001]{BEK00} Beristain, G., Edwards, S., \& Kwan, J.\ 2001, \apj,  551, 1037 
\bibitem[Binette et al.~1999]{Binette99} Binette, L., Cabrit, S., Raga, A., \& Cant{\' o}, J.\ 1999, \aap,  346, 260
\bibitem[Brugel, Boehm, \& Mannery(1981)]{1981ApJS...47..117B} Brugel, E.~W., B\"ohm, K.~H., \& Mannery, E.\ 1981, \apjs, 47, 117
\bibitem[Cohen \& Fuller 1985]{Cohen85} Cohen, M. \& Fuller, G.~A.\ 1985, \apj,  296, 620
\bibitem[Dougados, Cabrit, Lavalley, \& M{\' e}nard(2000)]{2000A&A...357L..61D} Dougados, C., Cabrit, S., Lavalley, C., \& M{\' e}nard, F.\ 2000, \aap, 357, L61
\bibitem[Eisl{\" o}ffel \& Mundt(1998)]{1998AJ....115.1554E} Eisl{\" o}ffel, J.~\& Mundt, R.\ 1998, \aj, 115, 1554
\bibitem[Eis\"offel, Mundt, Ray, \& Rodriguez(2000)]{2000prpl.conf..815E} Eis\"offel, J., Mundt, R., Ray, T.~P., \& Rodriguez, L.~F.\ 2000, Protostars and Planets IV (Book - Tucson: University of Arizona Press; eds Mannings, V., Boss, A.P., Russell, S.~S.), p.~815, 815
\bibitem[Folha \& Emerson 2001]{2001A&A...365...90F} Folha, D.~F.~M.~\& Emerson, J.~P.\ 2001, \aap,  365, 90
\bibitem[Garcia et al.~2001a]{Garcia01} Garcia, P.~J.~V., Ferreira, J., Cabrit, S., \& Binette L.\ 2001, \aap, 377, 589
%\bibitem[Garcia et al.~2001b]{Garcia01b} Garcia P.~J.~V., Cabrit S., Ferreira J., Binette L., 2001b, \aap,  377, 609
\bibitem[Goodson et al.~1999]{G1999} Goodson, A.~P., B{\" o}hm, K., \& Winglee, R.~M.\ 1999, \apj,  524, 142
\bibitem[Gullbring, Calvet, Muzerolle, \& Hartmann(2000)]{2000ApJ...544..927G} Gullbring, E., Calvet, N., Muzerolle, J., \& Hartmann, L.\ 2000, \apj, 544, 927
\bibitem[Hamann 1994]{H94} Hamann, F.\ 1994, \apjs, 93, 485
%\bibitem[Hamann et al.~1994]{HSCP94} Hamann F., Simon M., Carr J.~S., Prato L., 1994, \apj,  436, 292 
\bibitem[Hartigan et al.~1995]{HEG95} Hartigan, P., Edwards, S., \& Ghandour, L.\ 1995, \apj,  452, 736
\bibitem[Hartmann et al.~1994]{H94} Hartmann, L., Hewett, R., \& Calvet, N.\ 1994, \apj,  426, 669
%\bibitem[Hirth, Mundt, \& Solf(1994)]{1994A&A...285..929H} Hirth, G.~A., Mundt, R., \& Solf, J.\ 1994, \aap, 285, 929
\bibitem[Hirth, Mundt, \& Solf(1997)]{1997A&AS..126..437H} Hirth, G.~A., Mundt, R., \& Solf, J.\ 1997, \aaps, 126, 437
\bibitem[Kepner et al.~1993]{Kepner93} Kepner, J., Hartigan, P., Yang, C., \& Strom, S., 1993, \apjl,  415, L119
%\bibitem[Kitamura et al.~1996]{Kiramura96} Kitamura, Y., Kawabe, R., \& Saito, M.\ 1996, \apj,  457, 277
\bibitem[Kobayashi et al.(2000)]{2000SPIE.4008.1056K} Kobayashi, N.~et al.\ 2000, \procspie, 4008, 1056
\bibitem[Konigl \& Pudritz(2000)]{2000prpl.conf..759K} Konigl, A.~\& Pudritz, R.~E.\ 2000, Protostars and Planets IV (Book - Tucson: University of Arizona Press; eds Mannings, V., Boss, A.P., Russell, S.~S.), p.~759, 759
\bibitem[Lavalley et al.(1997)]{1997A&A...327..671L} Lavalley, C., Cabrit, S., Dougados, C., Ferruit, P., \& Bacon, R.\ 1997, \aap, 327, 671
\bibitem[Lavalley-Fouquet, Cabrit, \& Dougados(2000)]{2000A&A...356L..41L} Lavalley-Fouquet, C., Cabrit, S., \& Dougados, C.\ 2000, \aap, 356, L41
\bibitem[Lord(1992)]{lor92} Lord, S. D.\ 1992, NASA Technical Memor. 103957 
\bibitem[Muzerolle et al.~1998]{Muzerolle98a} Muzerolle, J., Calvet, N., \& Hartmann, L.\ 1998b, \apj,  492, 743
\bibitem[Muzerolle et al.~2001]{Muzerolle01} Muzerolle, J., Calvet, N., \& Hartmann, L.\ 2001, \apj,  550, 944
\bibitem[Muzerolle et al.~1998]{Muzerolle98b} Muzerolle, J., Hartmann, L., \& Calvet, N., 1998a, \aj,  116, 2965
\bibitem[Najita et al.~2000]{N2000} Najita, J.~R., Edwards, S., Basri, G., \& Carr J., 2000  Protostars and Planets IV (Book - Tucson: University of Arizona Press; eds Mannings, V., Boss, A.P., Russell, S.~S.), p.~457
\bibitem[Oudmaijer, Busfield, \& Drew(1997)]{1997MNRAS.291..797O} Oudmaijer, R.~D., Busfield, G., \& Drew, J.~E.\ 1997, \mnras, 291, 797
\bibitem[Safier 1993a]{Safier93a} Safier, P.~N.\ 1993, \apj,  408, 115
\bibitem[Shu et al.~2000]{Shu00} Shu, F.~H., Najita, J.~R., Shang, H., \& Li, Z.-Y., 2000\ , Protostars and Planets IV (Book - Tucson: University of Arizona Press; eds Mannings, V., Boss, A.P., Russell, S.~S.), p.~789
\bibitem[Takami et al.(2001)]{2001MNRAS.323..177T} Takami, M., Bailey, J., Gledhill, T.~M., Chrysostomou, A., \& Hough, J.~H.\ 2001, \mnras, 323, 177
\bibitem[Takami et al.~2002]{Takami02} Takami, M. et al. 2002, \apj, 566, 910
\bibitem[Tokunaga et al.(1998)]{1998SPIE.3354..512T} Tokunaga, A.~T.~et al.\ 1998, \procspie, 3354, 512
\bibitem[Wichmann et al.(1998)]{1998MNRAS.301L..39W} Wichmann, R., Bastian, U., Krautter, J., Jankovics, I., \& Rucinski, S.~M.\ 1998, \mnras, 301, L39
\end{thebibliography}