# Structure and Dynamics of the Accretion Process and Wind
in TW Hya ††thanks: Data presented herein were obtained at the W. M. Keck
Observatory, which is operated as a scientific partnership among the
California Institute of Technology, the University of California, and the
National Aeronautics and Space Administration. The Observatory was made
possible by the generous financial support of the W. M. Keck Foundation.
Infrared spectra were taken at the Gemini Observatory, which is operated by
the Association of Universities for Research in Astronomy, Inc., under a
cooperative agreement with the NSF on behalf of the Gemini partnership: the
National Science Foundation (United States), formerly the Science and
Technology Facilities Council (United Kingdom), the National Research Council
(Canada), CONICYT (Chile), the Australian Research Council (Australia),
Ministério da Ciência e Tecnologia (Brazil) and Ministerio de Ciencia,
Tecnología e Innovación Productiva (Argentina). This paper also includes
spectra gathered with the 6.5-meter Magellan Telescope/CLAY located at Las
Campanas Observatory, Chile. Additional spectra were obtained at the 1.5m
Tillinghast Telescope at the Fred Lawrence Whipple Observatory of the
Smithsonian Astrophysical Observatory.

A. K. Dupree, N. S. Brickhouse, S. R. Cranmer, P. Berlind22affiliation: Fred
L. Whipple Observatory, Amado, AZ , Jay Strader33affiliation: Michigan State
University, East Lansing, MI Harvard-Smithsonian Center for Astrophysics,
Cambridge, MA 02138, USA  Graeme H. Smith  University of California
Observatories/Lick Observatory, University of California,
Santa Cruz, CA 95064

###### Abstract

Time-domain spectroscopy of the classical accreting T Tauri star, TW Hya,
covering a decade and spanning the far UV to the near-infrared spectral
regions can identify the radiation sources, the atmospheric structure produced
by accretion, and properties of the stellar wind. On time scales from days to
years, substantial changes occur in emission line profiles and line strengths.
Our extensive time-domain spectroscopy suggests that the broad near-IR,
optical, and far-uv emission lines, centered on the star, originate in a
turbulent post-shock region and can undergo scattering by the overlying
stellar wind as well as some absorption from infalling material. Stable
absorption features appear in H$\alpha$, apparently caused by an accreting
column silhouetted in the stellar wind. Inflow of material onto the star is
revealed by the near-IR He I 10830Å line, and its free-fall velocity
correlates inversely with the strength of the post-shock emission, consistent
with a dipole accretion model. However, the predictions of hydrogen line
profiles based on accretion stream models are not well-matched by these
observations. Evidence of an accelerating warm to hot stellar wind is shown by
the near-IR He I line, and emission profiles of C II, C III, C IV, N V, and O
VI. The outflow of material changes substantially in both speed and opacity in
the yearly sampling of the near-IR He I line over a decade. Terminal outflow
velocities that range from 200 km s-1 to almost 400 km s-1 in He I appear to
be directly related to the amount of post-shock emission, giving evidence for
an accretion-driven stellar wind. Calculations of the emission from realistic
post-shock regions are needed.

stars: individual (TW Hydrae) - stars:pre-main sequence - stars: variables: T
Tauri, Herbig Ae/Be - stars: winds, outflows - Accretion, accretion disks

††slugcomment: May 5, 2014

##  1 Introduction

TW Hya (CD $-$34 7151, TWA 1, HIP 53911) remains arguably the closest
accreting T Tauri Star (Wichmann et al. 1998), and is oriented with its
rotation axis almost along our line of sight which sets the surrounding
accretion disk approximately in the plane of the sky (Krist et al. 2000; Qi et
al. 2004). These characteristics make TW Hya a subject of intensive study at
all wavelengths because it is bright and the accretion process is directly
accessible by this polar orientation (Donati et al. 2011; Johnstone et al.
2014).

TW Hya has been a frequent target of intensive study of optical emission lines
along with several other bright T Tauri stars: BP Tau (Gullbring et al. 1996);
DF Tau (Johns-Krull & Basri 1997); DQ Tau (Basri et al. 1997); DR Tau (Alencar
et al. 2001); and RW Aur (Alencar et al. 2005) among them. These studies
reveal the intrinsic variability of the optical lines. These stars have larger
inclinations than TW Hya, [20∘ for DR Tau (Schegerer et al. 2009); 39∘ for BP
Tau (Guilloteau et al. 2011); the star RW Aur A is inclined by 37∘; others are
not known]. The uniqueness of TW Hya resides in the imaging of the surrounding
disk in the infrared (Krist et al. 2000) and in the sub-millimeter range (Qi
et al. 2004) revealing its face-on orientation. Thus the interesting polar
region where accretion is ongoing can be viewed directly. This paper reports
spectral sequences not only from the optical region, but the near-infrared and
ultraviolet as well. The broad wavelength coverage spans accretion phenomena,
the presence of winds, and the intrinsic chromosphere/corona of the star.

TW Hya itself is of low mass. Estimates range from 0.4 – 0.8 M⊙ (Batalha et
al. 2002; Donati et al. 2011; Vacca & Sandell 2011). The spectral type is
believed to be close to a K7 dwarf (Alencar & Batalha 2002). A possible later
spectral type (M2.5V, Vacca & Sandell 2011) is controversial (Andrews et al.
2012; McClure et al. 2013). The inferred mass and radius for a M2.5V star lead
to lower free-fall accretion velocities and consequently shock temperatures
lower than measured directly (Brickhouse et al. 2010). This star has become a
fiducial object in the astrophysics of accretion and low mass star formation.

Currently, the ‘standard model’ of accretion suggests (Hartmann 1998) that
material from a surrounding disk is channeled by magnetic fields in an
‘accretion funnel’ towards the central star forming an accretion shock, a
small post-shock cooling zone, and a hot spot or ring on or near the stellar
surface. The source of optical emission is assigned to the accreting funnel
flows (Hartmann et al. 1994; Muzerolle et al. 2001; Kurosawa & Romanova 2013).
Alencar and Batalha (2002) in a very detailed analysis of 42 optical spectra
of TW Hya acquired mostly nightly over a 1.4 year period described the
behavior of several emission lines: H$\alpha$, H$\beta$, He I ($\lambda$5876)
and Na D. They noted both outflow and inflow signatures in these profiles and
documented a correlation between veiling corrected equivalent widths of major
lines and veiling. This suggests that increased veiling $-$ believed to arise
from the continuum and lines (Gahm et al. 2008) produced by the accretion hot
spot (or ring) $-$ is related to the increased equivalent width of optical
emission lines. The accretion rate and its variation can be inferred directly
from the Ne IX lines in the X-ray spectrum (Brickhouse et al. 2012). However,
recent optical spectroscopy simultaneous with X-ray measurements demonstrated
the progression of accreted material in the post-shock cooling zone through
the stellar atmosphere, producing optical emission, followed by the heating of
the photosphere, and subsequent enhancement of the corona (Dupree et al.
2012). This sequence challenges the ‘standard model’ which attributes optical
emission to the accretion funnels (Muzerolle et al. 2001, 2005; Natta et al.
2004). In addition, over this time, a stellar wind becomes established and
increases in strength over several days. X-ray spectroscopy of TW Hya and the
accompanying spectroscopic diagnostics of density and temperature revealed a
large post-shock volume in the corona (Brickhouse et al. 2010). This newly-
discovered material has roughly 300 times larger volume and 30 times more mass
than the accretion shock itself, and signals the presence of an accretion-fed
corona. It appears to be a large turbulent billowing structure in the corona.
This process may be similar to the recently identified brightenings in the
solar corona following the impact of fragments from an erupting solar filament
(Reale et al. 2013). Accreting material can supply large nearby magnetic
structures and also drive a stellar wind (Cranmer 2009).

While broad ultraviolet, far-UV, and optical emission lines from TW Hya are
well documented (Alencar & Batalha 2002; Herczeg et al. 2002; Dupree et al.
2005a; Ardila et al., 2013) their origin has remained elusive. Intensive study
of the X-ray spectrum of TW Hya with CHANDRA coupled with simultaneous optical
spectra strongly suggests that the source of its broad optical and X-ray
emission arises from the turbulent postshock cooling volume in the stellar
atmosphere (Brickhouse et al. 2010; Dupree et al. 2012), and not from
‘accretion funnels’ approaching the star. The broad emission can suffer some
absorption from the accretion stream and can also be substantially modified by
a stellar wind which appears to be driven by the accretion process. The
structure of this wind, its characteristics as well as the amount of mass loss
remain to be determined. These processes have substantial implications for
angular momentum loss from the star (Matt & Pudritz 2005; Matt et al. 2012;
Bouvier et al. 2014), for the presence of dust in the surrounding
circumstellar material (Alexander et al. 2005), and for an understanding of
the accretion process itself and its consequences.

To summarize, the recent measurements suggest a more complicated accretion
process than found in the ‘standard model’ described above. We envision a
magnetically channeled accretion flow from the circumstellar disk towards the
star. This is defined as an accretion funnel. This flow of plasma accelerates,
free-falling towards the star and forms a shock - an abrupt increase in plasma
temperature and density. The post-shock plasma, a turbulent medium, appears to
cool radiatively. The subsequent effects of the accretion shock are many: a
turbulent cooling region that produces the broad ultraviolet emission, the
broad optical emission lines, and the broad near-IR helium emission; a heated
spot or ring in the stellar photosphere that creates the veiling or filling-in
of photospheric lines; a large heated volume in the stellar corona detected in
O VII; and acceleration of a stellar wind detected as absorption in H$\alpha$,
He I and ultraviolet line profiles.

In this paper, many new sequences of optical, near infrared, and far
ultraviolet spectra are presented to infer the characteristics of the
accretion process and wind on several time scales. The spectra allow
identification of new structures in the stellar atmosphere, give insight into
the accretion parameters, and probe the wind from the star. Spectral
observations in the far-UV, optical, and near-IR regions are described in
Section 2; the source of the H$\alpha$ line is discussed in Section 3\.
Spectral indications of the accretion funnels themselves are presented in
Section 4. The H$\delta$ transition is compared to predictions of current MHD
models in Section 5. The variable warm wind111Here we take a warm wind to have
a temperature of $\sim$15,000K indicative of the formation region of He I; a
hot wind is considered to have a temperature of $\sim$105 K or higher. A
stellar wind, if like the solar wind, exhibits a progression in temperature
from ‘warm’ to ‘hot’. extending to greater heights in the chromosphere than
H-$\alpha$ is shown by the near-infrared transition of He I (Section 6).
Section 7 presents evidence inferred from the far ultraviolet spectra bearing
on the variable post-shock conditions. Section 8 evaluates a previous claim of
wind presence and characteristics. Discussion and Conclusions occur in Section
9.

##  2 Observations

A variety of spectroscopic observations (Table 1) has been acquired to assess
the structure and dynamics of the atmosphere of TW Hya. Optical spectra were
obtained during many observing runs at the Magellan-CLAY 6.5m telescope at Las
Campanas Observatory, Chile. MIKE, the double-echelle spectrograph (Bernstein
et al. 2003) was used with a 0.75″$\times$5″ slit that yields a 2-pixel
resolution element of $\lambda/\Delta\lambda\sim$30,000 on the red side
($\lambda\lambda$4900–9300) which contains the H$\alpha$ line and
$\lambda/\Delta\lambda\sim$37,000 on the blue side ($\lambda\lambda$3200–5000)
used for the higher Balmer lines. The IDL pipeline developed by S. Burles, R.
Bernstein, and J. X. Prochaska222See http://web.mit.edu/~burles/www/MIKE/ was
used to extract the spectra, and IRAF software333See iraf.noao.edu was
employed in the analysis. Additional optical spectra were taken at the FLWO
1.5-m telescope with the Tillinghast Reflector Echelle Spectrograph (TRES).
TRES is a temperature-controlled, fiber-fed instrument with a resolving power
R$\sim$44,000. The spectra were reduced with a custom IDL pipeline.444See
www.sao.arizona.edu/FLWO/60/tres.html

Near-infrared spectra were obtained with PHOENIX at Gemini-S during two
classical observing runs. The PHOENIX setup consisted of a slit width of 4
pixels yielding a spectral resolution of $\sim$50,000. The order-sorting
filter J9232 was selected which spans 1.077–1.089$\mu$m and allows access to
the He I line at 1.083$\mu$m. Standard procedures were followed by acquiring
target spectra using a nodding mode (A–B) with a spatial separation of 5
arcsec. Because standard comparison lamps have a sparse wavelength pattern in
this near-infrared region, we observed a bright K giant containing many
securely identified narrow photospheric absorption lines in order to determine
the wavelength scale. Other near-infrared spectra were taken with NIRSPEC
(McLean et al. 1998, 2000) on the Keck II telescope. The echelle cross-
dispersed mode of NIRSPEC with the NIRSPEC-1 order-sorting filter was selected
and the slit of 0.42″$\times$12″ gives a nominal spectral resolution of
23,600. The long-wavelength blocking filter was not used in order to minimize
unwanted fringing. Internal flat-field lamps, NeArKr arcs, dark frames, and K0
giants were used as calibration exposures. Spectra were taken in a standard
nodding mode. Data reduction was carried out with the REDSPEC package
specifically written for NIRSPEC (McLean et al. 2003).

Far-UV spectra of TW Hya from FUSE were taken from the MAST archive and the
extraction procedure was fully described elsewhere (Dupree et al. 2005a). Some
HST spectra of TW Hya obtained with STIS were also taken from the MAST archive
at STScI, and from the CoolCAT UV spectral catalogue (Ayres 2005).555See
casa.colorado.edu/~ayres/CoolCAT/. Other HST spectra came directly from the
MAST archive at STScI. Details of these spectra are also given in Table 1.

##  3 Identifying the Source of H$\alpha$

The H$\alpha$ profile is useful for probing the accretion process. A broad
H$\alpha$ line has long been considered a signal of ongoing accretion in young
stars (Bertout 1989; Hartmann et al. 1994; Alencar & Batalha 2002), and the
cause of the broadening has frequently been ascribed to ‘turbulence’ (Alencar
& Basri 2000). In this section, we present the most extensive study of TW Hya
in the optical region. The H$\alpha$ profile was measured on 41 nights over a
decade (Figure 1), with anywhere from 2 to 300 spectra per night. Most of the
observations were taken on sequential days which enables the identification of
dynamic events in a stellar atmosphere. As has been noted almost two decades
ago (Johns-Krull & Basri 1997), snapshot spectra do not provide a
characteristic representation of a complex accreting system, and time
sequences are invaluable. The H$\alpha$ profiles in TW Hya show dramatic
changes in the strength of the emission and the presence of absorption
features at negative (outflowing) velocities.666Water vapor produces an
absorption feature on the positive velocity side near $+$50 km s-1. The
signature of a wind in the H$\alpha$ profile has been noted previously
(Hartmann 1982; Calvet & Hartmann 1992; Gullbring et al. 1996; Alencar & Basri
2000; Alencar & Batalha 2002) but only recently has the hourly development of
wind absorption became evident (Dupree et al. 2012). The shape of the emission
profile indicates material motions of gas at chromospheric temperatures.
Comparison of profiles in Figure 1 and evaluation of the flux at negative and
positive velocities demonstrates the presence of roughly symmetric H$\alpha$
profiles (to within 10%) centered on TW Hya for 42% of the spectra and a
positive flux enhancement (‘red enhanced’) for 52% of the spectra. The full
width at half maximum (FWHM) of the lines rangew from 200 to 300 km s-1. This
positional symmetry is consistent with formation on the star itself. When the
profiles appear to be ‘red-enhanced’, this is caused by substantial absorption
and scattering (on the “blue-side”) at negative velocities by outflowing
material forming the stellar wind. A particularly instructive sequence
occurred in 2007 Feb (Dupree et al. 2012). At that time, the H$\alpha$ profile
was roughly symmetric on the first night of observation, then the wind opacity
systematically increased over the following 3 nights.

These observations demonstrate that the H$\alpha$ profile from TW Hya is
inherently broad and symmetrically positioned on the stellar velocity, but can
be substantially absorbed by the stellar wind. In addition, changes in the
long-wavelength side of the profile (cf. Figure 1) can result from a
combination of enhanced post-shock emission and changing absorption by the
accretion stream. The pole of TW Hya is observed directly. It is here that
accretion occurs most of the time (Donati et al. 2011), and the shock and
post-shock cooling region can be seen. The source of this broad emission has
been long debated. With the assumption that emission arises from the accretion
stream, models tend to predict pointed profiles from a star observed at low
inclination. Frequently absorption on the red side that results from the
infalling material emerges from the calculations (Muzerolle et al. 2001). This
makes H$\alpha$ appear shifted to shorter wavelengths in contrast to most of
these observations. Authors have noted that the magnetospheric accretion
models do not adequately represent spectral profiles found in several T Tauri
stars (e.g. Alencar & Basri 2000). Changes in damping parameters and the
introduction of Stark broadening, can manage to broaden the H$\alpha$ profile
but the observed profiles in TW Hya have flatter peaks and are frequently wind
absorbed in contrast to the calculated profiles (Figure 2).

During many nights the observed H$\alpha$ profile remains centered on the
stellar radial velocity and appears symmetric without a distinct absorption
feature. As previously suggested (Dupree et al. 2012), the broad H$\alpha$
lines can result naturally from turbulent conditions in a post-shock cooling
volume. This identification derived from the time study of an X-Ray accretion
event which was followed by sequential changes in emission line profiles and
fluxes, veiling, and subsequent coronal X-Ray enhancement. The breadth of the
H$\alpha$ line is slightly larger ($\sim$200$-$275 km s-1), but comparable to
that of Ne IX X-ray lines arising in the shock (126 to 229 km s-1) while
slightly less than the full width of the symmetrized far-ultraviolet lines (C
III and O VI) observed to be $\sim$ 325 km s-1 in TW Hya (Dupree et al.
2005a). Moreover, the amount of material in the post-shock region inferred
from the X-ray diagnostics (Brickhouse et al. 2010) as well as near-infrared
measurements (Vacca & Sandell 2011) exceeds by a few orders of magnitude the
material expected in an accretion column (cf. Dupree et al. 2012). These facts
strongly suggest that the symmetric H$\alpha$ emission line arises
predominantly from the turbulent post-shock cooling region. It has been known
for a long time that models of a stellar chromosphere alone can not produce
strong hydrogen emission lines (Calvet et al. 1984). A symmetric H$\alpha$
profile centered on the stellar radial velocity is expected from TW Hya
because the postshock region is in our full view. When a wind from the star is
present, the short wavelength side of the profile is weakened, and the
resulting profile is stronger on the long wavelength side of the line as
frequently appears in TW Hya. Absorption by infalling material can affect the
profile as well.

Systems with different inclinations might be expected to produce similar
turbulently broadened profiles if the post-shock region is not obscured by a
disk. In fact, broad H$\alpha$ has long been known (Bertout 1989; White &
Basri 2003) to signal the presence of a classical accreting T Tauri star.
Differences in the profile may arise from wind scattering which would be
determined by the line-of-sight component of the expansion velocity. But as is
evident from Figure 1, where the line-of-sight reflects the maximum outflowing
velocity from TW Hya, the wind velocity and opacity can change substantially
with time. The He I profile, discussed in Section 6 offers an interesting test
of the effects of inclination.

##  4 Detection of Shadowing by Accretion Funnels

The H$\alpha$ profiles also reveal structures that occur in the stellar wind
and their behavior becomes apparent in the spectral sequences lasting several
days. Alencar and Batalha (2002) remarked on a broad (FWHM $\sim$125 km s-1)
absorption feature in H$\alpha$ that appeared in nightly spectra suggesting
variation in wind opacity. Time domain spectroscopy presented here confirms
the changes in the wind, and also reveals another feature. Profiles in Figure
1 show an absorption feature that occurs at $\sim$$-$50 to $-$100 km s-1. This
feature can change in strength and velocity from night to night, but is
constant in velocity over a single night. Particularly instructive are the
nights following 2004 Apr 27 (Figure 3) where the absorption feature appears
at $-$80 km s-1, vanishes on Apr 28, re-appears on Apr 29 at $-$150 km s-1,
and on Apr 30 at $-$100 km s-1. Similar behavior occurred in 2007 Feb and 2009
Jun (Figure 1) and also in 2006 (Figure 4). These absorption features in the
wind remain at constant velocity during the night. The three-night sequence
from 2006 April is presented in gray scale representation in Figure 5. The
first night shows a narrow ($\sim$50 km s-1) absorption feature at $-$125 km
s-1 at constant velocity; this feature becomes weaker early in night 2.
However, during the second night of observations, the absorption at $-$50 km
s-1 at the start of the night is replaced by an absorption feature at $-200$
km s-1 by the end of the night. The right panel of Figure 5 which contains a
line plot of the 8 hours of observation during night 2 illustrates this jump
quite clearly. In addition, the variable infall of material at high velocities
can be seen in the profile. The short rotation period of TW Hya (3.57 days;
Huélamo et al. 2008) causes different viewing angles to occur during the 3
successive nights of our observations. The narrow absorptions appear constant
in velocity suggesting they are located at high latitudes on the star. If the
features moved outward within the wind, the speeds of $\sim$150 km s-1 would
allow material to cover many ($\sim$6) stellar radii over an 8-hour span of
observations and this motion would also be visible in the line profiles. But
such a shift is not observed.

Such static behavior differs from high-lying stellar prominences found in
rapidly rotating stars such as AB Dor (Collier Cameron & Robinson 1989) and in
rapidly rotating T Tauri stars (Oliveira et al. 2000; Skelly et al. 2008;
Günther et al. 2013) where the absorption feature moves rapidly in velocity
from negative to positive (or vice versa) as the feature traverses the stellar
disk. The absorption features in TW Hya are stable in velocity for hours, and
are not at all similar to the moving Discrete Absorption Components (DACs)
observed in the winds of O stars (cf. Howarth et al. 1995) which generally are
very broad ($\sim$400 km s-1) and move at thousands of km s-1 through the
atmosphere (cf. Kaper et al. 1999); only absorption features in one Wolf-Rayet
star (HD 50896) are broad, but static (Massa et al. 1995).

The static nature of these relatively narrow absorptions in TW Hya suggests
they arise in a stable cool structure located in the wind itself that is
silhouetted in the hydrogen line profile. These could be the signature of the
accreting columns crossing the path of the wind. In an accelerating wind,
features located at $-$50 to $-$150 km s-1, would exist at about 1.2-1.3 R⋆
(Dupree et al. 2008), judging from the chromospheric helium line asymmetries,
and above the shocked material produced by accretion (Sacco et al. 2010). The
strength of the absorptions, from 0.4 to 0.6 of the local continuum provided
by the H$\alpha$ emission, indicates an optical depth, $\tau$ of $\sim$0.7.
Here we take $\tau=\kappa\rho H$ where $\kappa(T,n_{2}/n_{1})$ is the line
absorption coefficient (incorporating a dependence on temperature and the
level population ratio of hydrogen), $\rho$ is the hydrogen density, and $H$,
the thickness of the line-forming region. The temperature and density
structure of the ‘funnel flows’ is uncertain. However, the preshock electron
density was determined (Brickhouse et al. 2010) to be 5.8$\times$1011 cm-3
which implies a hydrogen density of 5$\times$1011 cm-3 for a plasma with 10%
helium. For a temperature T $\sim$7000K, the observed broadening and our non-
LTE calculations777The PANDORA code (Avrett & Loeser 2008) was used in a
spherical expanding semi-empirical model constructed to match wind scattering
profiles of H$\alpha$, He I, and the chromospheric density inferred from the
ultraviolet C II lines. for chromospheric hydrogen suggest that the n2/n1
level populations of the H$\alpha$ line range from 6.4$\times$10-7 to
1.8$\times$10-8 in this temperature and density regime (Dupree et al. 2008).
The scale length to produce the narrow absorption is small, less than 0.01 R⋆.
Another model of the accreting stream (Hartmann et al. 1994) gives T = 7000K
and the level 2 population of hydrogen as 8 $\times$105 cm-3 near the star,
such that the inferred thickness is much smaller, $\sim$100 km. In either
case, formation in a narrow accretion funnel appears plausible to account for
these absorption features.

Chromospheres, winds, and wind structures produce asymmetric lines and
absorption features that may be broad (such as a central reversal) or narrow
as discussed here. The uniqueness of these narrow ‘notches’ relates to their
constant velocity in the H$\alpha$ line profile. A blocking structure located
in the wind could be seen at all angles where it is observed against a region
producing emission. Many studies of emission from T Tauri stars have been
published but very few have the time sampling of multiple sequential spectra
that are presented here.

##  5 The H$\delta$ Line

Sophisticated models of magnetospheric accretion processes have also been
constructed and used to predict hydrogen line profiles. A sequence of Balmer
lines calculated by Kurosawa & Romanova (2013) suggested that the H$\delta$
transition at 4101Å is an optimum line to reveal persistent red-shifted sub-
continuum absorption resulting from the inflowing accretion stream in the
‘standard’ model. Absorption is expected when the hotspot is in the line of
sight of the observer as is the accretion stream. TW Hya is a good test of
these assumptions. Because of its pole-on orientation, the accretion stream
and hot spot are in the line of sight. Our MIKE spectra capture the pivotal
H$\delta$ line, as well as the complete Balmer series. Figure 6 shows four
lines of the Balmer series (H$\alpha$, H$\beta$, H$\gamma$, H$\delta$) taken
over 4 sequential nights in 2007 (Dupree et al. 2012). During this period the
H$\alpha$ line displays a slight inflow asymmetry (blue side stronger than red
side) early on the first night, followed by the subsequent onset of a wind.
H$\beta$ and H$\gamma$ exhibit an absorption feature at $\sim+$50 km s-1 which
signals enhanced inflow following the X-ray accretion event (Dupree et al.
2012). However the predicted sub-continuum absorption signalling infalling
material is not evident in the H$\delta$ line. The period of radial velocity
variation, namely 3.57 days (Huélamo et al. 2008) could cause a change in the
absorption strength if there were substantial modulation in the orientation of
the hot spot and the accretion column. The lack of an absorption feature in
the H$\delta$ profile suggests that the H$\delta$ emission does not arise from
infalling material, but must originate from another region. Another sequence
of H$\delta$ profiles over 4 consecutive nights is shown in Figure 7. The
first night of the sequence shows enhanced emission with shallow broad
absorption between $+$100 to $+$200 km s-1, inferred from the line asymmetry.
This pattern may be the signal of an accretion event. However, subsequent
nights (July 14$-$16) reveal a symmetric profile, indicating the infall has
lessened, and the symmetry and FHWM ($\sim$150 km s-1) are consistent with
formation in a turbulent post-shock cooling volume. Variability of the
accretion rate is also documented by the X-ray diagnostics (Brickhouse et al.
2012) which have shown a five-fold reduction in the rate from exposures
separated by 2.7 days. Other studies have measured H$\delta$ in T Tauri stars
(Edwards et al. 1994; Petrov et al. 1996; Alencar & Basri 2000) and the
profiles in the majority of stars do not resemble the models presented in
Kurosawa & Romanova (2013).

##  6 Variable Warm Wind in Near IR Helium

The broad near-IR line of He I at 10830Å, because of its width (FWHM $\sim$200
km s-1), also appears formed in the post-shock material. This feature has
become a useful probe of the wind and infall environment of accreting T Tauri
stars (Dupree 2003; Edwards et al. 2003; Dupree et al. 2005a; Edwards et al.
2006). For TW Hya, we have acquired near-IR spectra spanning 8 years including
sequences of consecutive nights. A comparison of the spectra obtained during
this 8 year span is shown in Figure 8. Over this long time scale, practically
every characteristic of the P Cygni profiles changes: the emission level, the
wind speed and opacity, and the appearance (or not) of inflowing material. The
emission strength varies by more than a factor of 2; the terminal velocity of
the wind varies from $-$200 to $-$315 km s-1; the wind opacity changes over
the absorption profile; the infall terminal velocity varies from +270 to +363
km s-1 as does the opacity of the infalling material. Speaking precisely, an
outflow of material is observed in the near-infrared helium profile of TW Hya.
The escape velocity from the photosphere is 530 km s-1 for TW Hya assuming
M=0.8M⊙ and R=1.1R⊙ (Donati et al. 2011). At a distance of 1 stellar radius
above the surface, this value decreases to 400 km s-1, which is comparable to
the observed typical terminal velocities. It is highly likely then that the
outflows observed in the near-IR helium line and the sequential acceleration
with temperature displayed by the UV lines in Dupree et al. (2005a) form a
true stellar wind.888Other studies (Batalha et al. 2002) suggest M=0.7M⊙ and
R=0.8R⊙ which would indicate a comparable escape velocity (410 km s-1) at 1
stellar radius above the photosphere.

The time-domain spectroscopy revealing these changes suggests that all T Tauri
objects undoubtedly exhibit similar variability that snapshot studies (Edwards
et al. 2006) can not discern. The extent of the variation also blurs the
search for correlations of wind parameters with geometric inclination of the
star or disk (Appenzeller & Bertout 2013).

The helium absorption profile is broad, signalling its origin in a stellar
wind, through its classic P Cygni shape. There is no sign of a narrow low-
velocity ($\sim$10 to 100 km s-1) feature such as identified in several
spectra of other T Tauri stars (Edwards et al. 2006) and conjectured to
originate from the circumstellar disk. In the case of TW Hya, H2 and [O I]
emission arise from the disk (Herczeg et al. 2002, Pascucci et al. 2011) since
their velocities coincide with that of the star itself. The infrared [Ne II]
transition exhibits a low velocity ($-$5 km s-1) outflow which suggests the
presence of a cool (1000 K) photoevaporative molecular wind (Rigliaco et al.
2013). It is hard to see how helium which is more highly excited would be
produced in such a wind.

Several sequences of the He I line over consecutive nights are shown in Figure
9. At these times minimal variation of the line profiles occurred on a short
time scale: a small change in the emission in 2009; changes in the inflow
opacity in 2010. During 2009 and 2010 the wind absorption had constant high
optical depth (a flat-bottomed profile) over a 50 to 100 km s-1 span, and
subtle changes in opacity would not be observable. In contrast, the H$\alpha$
profile over sequential nights reveals changes in wind opacity and absorption
troughs (cf. Fig. 1 A-E, I, K, L). In 2009, the absorption feature at $-$50 km
s-1 is also found in the H$\alpha$ profile from those nights (see Figure 1).
At that time, the H$\alpha$ line revealed wind scattering that extends to
$-$100 km s-1, and the helium wind extended to $-$200 km s-1, an indication of
accelerated expansion.

The emission strength exhibits substantial change. Of course, the amount of
emission depends on the intrinsic line strength which is then modified by the
wind scattering and some absorption from infalling material. Close inspection
of the profiles suggests that the emission itself is indeed changing. On one
occasion, 2002 May 20, the profile displays emission above the continuum level
extending to both $-$300 and $+$300 km s-1 indicating a strong and broad
intrinsic emission component.

The inflow signature in helium is a particularly useful diagnostic of the
accretion process. In a standard magnetospheric accretion model (Brickhouse et
al. 2012), the terminal velocity of the infalling material, determined by the
free-fall distance from the circumstellar disk, controls the temperature of
the shock. X-ray spectra indicate that the shock temperature, measured
directly from the ratio of forbidden plus intercombination to resonance line
fluxes of Ne IX, varied by a factor of 1.6 over a span of 13 days between
CHANDRA pointings (Brickhouse et al. 2012). The shock temperatures vary
between 1.9MK and 3.1MK. In this standard model, the post-shock temperature is
given by $T_{post}\sim 3mv_{ff}^{2}/16k$, where $v_{ff}$ is the free-fall
velocity, $k$ is the Boltzmann constant and $m$ is the mean atomic mass
(Brickhouse et al. 2010). The measured shock temperatures from 3 CHANDRA
pointings in 2007 suggest that terminal free-fall velocities from 380 to 485
km s-1 occurred.

In the model (Brickhouse et al. 2012), when the free-fall velocity is low, the
disk supplying the accreting material originates closer in to the star. The
filling factor on the stellar surface increases which implies an increase in
the mass accretion rate, and also in the material in the post-shock cooling
zone. A range of parameters for this model is shown in Figure 10 (left panel).
This scenario is confirmed by the behavior of the helium line. The amount of
emission in the He I line appears inversely linked to the terminal velocity of
the inflowing material (Figure 10, right panel). Lower values of the inflow
terminal velocity lead to much stronger emission arising in the post-shock
region, which is consistent with the higher mass accretion rate as predicted
in the dipole model. The helium absorption indicates lower velocities than the
terminal velocities of the models suggesting that the absorption arises from a
volume above the accretion shock. It is also quite likely that the neutral
helium is ionized by the X-ray emission from the shock as the helium
approaches in the pre-shock stream (Lamzin 1999; Gregory et al. 2007;
Brickhouse et al. 2010), and as a result does not reach a terminal free-fall
velocity.

It is interesting that the helium line gives evidence of subcontinuum
absorption due to inflow, when the hydrogen lines generally do not show such
strong signatures of inflowing material. On 2007 March 1, we obtained (almost)
simultaneous spectra of helium (Figure 8) and the hydrogen lines (shown as
night 4 in Figure 6). The reversal of the negative velocity side of the Balmer
lines on to the positive velocity side (Figure 6) indicates weak broad
absorption near $+$200 km s-1 on the H$\beta$, H$\gamma$, and H$\delta$
transitions. However the helium line displays substantial absorption extending
to $+$325 km s-1. The temperature structure along the accretion column is not
known but it appears likely that helium forms at higher temperatures where the
hydrogen is ionized, and these temperatures are associated with higher
infalling velocities. This could create a higher opacity in helium with the
ability to probe a different, potentially more extended accreting region than
represented by the hydrogen lines.

Because stronger emission signals enhanced accretion of material, this
material might cause an increase in both the wind opacity and wind speed if
the accretion contributes to wind acceleration. There is a hint of this in
nine out of the ten near-IR spectra of He I (Figure 11). The one outlier with
the fastest outflow speed belongs to the 2007 March 1 observation discussed
above. Not only does this profile exhibit the highest outflow velocity, but
also the largest subcontinuum ‘blue absorption’ of all ten nights of
observations. It is difficult to quantify this further because the intrinsic
profile of the helium line is unknown. The profile exhibits absorption on the
‘red’ side caused by accretion, and substantial scattering on the ‘blue’ side
in the stellar wind. Figure 11 shows that increased emission in He I 10830Å
may be linked directly to a faster outflowing terminal velocity, giving
evidence for an accretion-driven stellar wind.

##  7 Variable Post-Shock Conditions from UV lines

The UV and far-UV lines also exhibit changes in flux as evidenced by HST and
FUSE spectra. The expanding atmosphere of TW Hya is revealed by the asymmetric
profiles of the major resonance lines (Dupree et al. 2005a), but substantial
changes in the profiles are evident at different times. Figure 12 (left panel)
contains resonance line profiles of C IV (1548Å) and N V (1239Å) measured with
HST/STIS on 8 May 2000. The line profiles are strikingly similar to one
another, confirming formation in the same region of the atmosphere. Ardila et
al. (2013) also remarked recently on this C IV/N V similarity from a much
larger sample of accreting stars. The temperatures of formation of C IV and N
V in a collisionally ionized plasma are similar at 1–2$\times$105K. The
profiles exhibit both extended emission to positive velocities, and the
characteristic sharp cutoff caused by wind absorption at negative velocities.
More wind opacity is evident in the C IV profile than in N V. This is to be
expected since the strength of absorption is proportional to the value of
$gf\lambda\times N_{i}/N_{H}$ where $N_{i}$ is taken to be $N_{C}$, $N_{O}$,
or $N_{N}$ giving the appropriate elemental abundance. The value of
$gf\lambda\times N_{C}/N_{H}$ for carbon is larger than the similar quantity
for the N V transition. Note the extended wing at positive velocities on this
date giving a half-maximum value of $+$200 km s-1. Contrast these profiles
with the resonance lines of O VI and C III on 20 Feb 2003 (Figure 12, right
panel) measured with FUSE. Both of these lines are also similar to one
another, and their formation temperatures are typically
8$\times$104–3$\times$105K in equilibrium conditions. The C III profile
exhibits more wind opacity than O VI which again can be understood from the
difference in the quantity $gf\lambda\times N_{C}/N_{H}$ which is larger for C
III than for O VI by a factor of 2.8. The relative amount of absorption in the
wind is in harmony with atomic physics. Thus, based on these spectra, we
conclude that the UV and far-UV lines are produced in the same plasma region,
but conditions in this region change dramatically from time to time.

Intercomparison of two of the ultraviolet profiles (Figure 13) reveals the
meaningful differences. Excess emission on the positive velocity side of the
profile in 2000 does not occur in the 2003 observation. The half-maximum value
of the O VI profile in 2003 extends only to $+$125 km s-1, whereas C IV
extended to $+$200 km s-1 in 2000. Additionally there is greater wind
scattering on the negative velocity side of the profile in C IV as compared to
O VI. While changes in the wind opacity between the two observations are
unknown, the atomic physics of these lines must be a first consideration. C IV
possesses a 17% larger atomic parameter ($gf\lambda\times N_{i}/N_{H}$) than O
VI (where $N_{i}$=$N_{C}$ or $N_{O}$) which contributes to its greater
opacity. We also do not know the distribution of material with temperature in
the wind which can affect both the ionization equilibrium and the formation of
the lines.

Variation of the C II and C IV emission on shorter time scales can be found in
Figure 14 from HST/STIS spectra. Substantial changes are apparent on a time
scale of years between the HST measurements, and also over a time scale of 6
days which separates the 2010 Jan 29 and 2010 Feb 4 observations of C IV. The
C IV emission appears to strengthen with C II, however the 4 simultaneous
measurements show scatter. The C II line is not broadened as much as the C IV
emission, suggesting that the contribution from the turbulent post-shock
cooling region is less. The positive velocity side of the C II emission
($\lambda$1335.7) resembles the positive velocity side of the H$\delta$ line
(cf. Figure 7) indicating that both may be formed in a region distinct from
the higher temperature species (C III, C IV, etc.).

A closer look at the variations in the high temperature emissions is provided
by the sequence of individual segments of the FUSE O VI spectra. Under
collisionally-dominated equilibrium conditions, O VI is formed at a
temperature of 3$\times$105K. However, the O VI line continues to be formed in
hotter coronal plasmas due to the extended wing of Li-like ions in the
ionization balance at high temperatures. O VI has been used consistently as a
diagnostic of the dynamical properties of outflows in the solar corona (Kohl
et al. 2006) and X-ray emitting shocks in OB-star winds (Lehner et al. 2003).
Evidence of the dramatic changes in emission from TW Hya can be found from the
individual C III and O VI profiles assembled over a 32 hour span and shown in
Figure 15. Each segment represents an integration time ranging from 2600 to
3700 s. The impression that C III varies more than O VI is quantitatively
confirmed by an extraction of the line fluxes. C III varies by a factor of
$\sim$1.8 and O VI by a factor of $\sim$1.5. Evaluating the negative and
positive velocity sides of each line shows that the fluxes are not correlated
between the sides. However the uncertainty in flux on the negative velocity
side is larger because the total counts are lower. And the total counts in the
O VI line are a factor of 4 or more than in the C III line, making the O VI
parameters a more secure measurement. The relatively constant profile of the
outflowing wind absorption at negative velocities contrasts dramatically with
the changing emission at positive velocities during this time. Such profile
variation provides additional evidence that the emission arises from the
turbulent post-shock accretion volume, and the the line is scattered in the
hot wind which is present during these observations. Similar behavior in He I
5876Å was found in another accreting T Tauri star, RU Lupi (Gahm et al. 2013).

Fluxes taken from the individual spectra are shown in Figure 16 where in the
first pointing the C III ($\lambda$977) line is stronger than O VI
($\lambda$1031.91); this behavior is reversed in the second pointing about one
day later. Here we have selected segments where the SiC detector (containing
the C III line) is well-aligned with the LiF detector (containing the O VI
line).999The FUSE telescope contained 4 coaligned optical channels. Pointing
on a target was maintained by the fine error sensor which viewed one of the
channels. Because thermal changes on orbit caused rotation of the mirrors,
this could lead to misalignment of the optical channels and a target could
drift in and out of the aperture in the channels not used for guiding. This
could affect the measurement of the C III line since it occurs in a channel
not used for guiding. In our analysis, segments for C III were selected when
the channels were coaligned. However, the C III atom has a metastable level
whose population is dependent on density which could contribute to the
behavior of C III. And of course, the post-shock cooling region can change as
well. During the second pointing when both lines could be measured
simultaneously over $\sim$5 hours, the fluxes are correlated with a linear
correlation coefficient of 0.9.

##  8 Remarks concerning a Hot Wind

The presence of a hot wind from TW Hya has been questioned by Johns-Krull and
Herczeg (2007), hereafter JKH. These authors offer several arguments
purporting to show that the wind is not hot. We argue below why we believe
that their concerns are not valid, and discuss four relevant issues.

(1) As evidence for outflow, JKH require sub-continuum absorption in a line
profile to be present at ‘blue-shifted’ wavelengths. While that phenomenon is
well-understood from characteristic P Cygni line profiles, the signature of
outflowing material can be revealed in other ways. A shift of the apparent
wavelength of the centroid of an emission or absorption line can reveal mass
motions. Mass motions can also be signaled by line asymmetries. Hummer and
Rybicki (1968) first pointed out that apparently redshifted line profiles can
arise in a differentially expanding atmosphere; in fact these have been widely
observed in luminous stars (cf. Mallik 1986, Robinson & Carpenter 1995, Dupree
et al. 2005b, Lobel & Dupree 2001) and in the Sun (Rutten & Uitenbroek 1991).
The assertion by JKH that sub-continuum absorption must be present to identify
outflow neglects both line-transfer physics and stellar observations
themselves.

(2) In an attempt to demonstrate the absence of absorption in resonance
doublets, JKH arbitrarily scale and overlay the long wavelength member of the
C IV doublet on to the short wavelength member of the doublet by forcing both
the line peaks and blue wings to match. They claim that signs of absorption
are missing in the red wing of the 1548Å line that would be caused by the wind
absorption in the 1550Å line. Such arbitrary scaling is not justified. Atomic
physics, supported by laboratory measurements, specifies the wavelength
separation between these lines as well as the ratio of the line oscillator
strengths. In an effectively optically thin doublet such as C IV, the line
flux ratio ($\lambda$1548/$\lambda$1550) is 2. Since it is preferable not to
amplify the errors in the weaker line of the doublet, here we place the
stronger (1548Å) line onto the weaker (1550Å) transition. The original HST
data used by JKH, namely O58D0130, has been superseded by a more optimal
spectrum reduction, namely CoolCAT (Ayres 2010) and we also consider that
reduction (see Table 1). In Figure 17, we show the CoolCAT reprocessed profile
as well as the original spectrum used by JKH.101010More recent papers (Yang et
al. 2012) than Herczeg et al. (2002) estimating the UV line fluxes from TW Hya
using revised calibrations report a substantially increased value of the C IV
flux by a factor of $\sim$5.

When the C IV 1548Å line is properly overlaid on the 1550Å line by shifting it
according to the wavelength separation (2.577Å, Griesmann & Kling 2000), and
scaling the flux of 1548Å by 2, two facts are immediately apparent from either
data reduction: (1) the wind opacity in the resonance line 1548Å, indicated by
the negative velocity side of the profile, is greater than in the subordinate
line 1550Å, as expected from atomic physics, and (2) excess wind absorption
does indeed occur in the long wavelength wing of 1548Å caused by wind
absorption from the 1550Å line. Consideration of the errors in the measured
fluxes at individual wavelengths in the HST/STIS spectrum indicates that the
separation visible in Figure 17 between the 1550Å line and the shifted scaled
1548Å line ranges from $\sim$2.5 to 3.2$\sigma$ in the velocity region from
$+$175 to $+$350 km s-1. Moreover, in this region the dashed line in Figure 17
systematically lies below the solid curve. Thus the effects of hot wind
absorption are clearly evident in the C IV profiles.

(3) Trying to demonstrate the absence of wind absorption by the O VI 1037.61Å
transition affecting the C II 1037.02Å line, JKH resort to an ad hoc
construction. In this case, both the activity of TW Hya and atomic physics
undermine their arguments. The O VI 1037.61Å line, has a neighboring C II line
separated by $-$172 km s-1 which places the C II line in the velocity range of
scattering produced by a hot O VI wind. To attempt to show that the C II line
is not weakened by scattering, JKH predict the intrinsic strength of a blended
O VI and C II feature. By taking the C II line profile at 1335.71Å observed in
2000 with HST/STIS, these authors scale the profile to mimic the C II
transition at 1037.02Å, and add it to the line profile of O VI 1031.91Å
observed in 2003. The resulting scaled combination is then compared to the
blended O VI(1037.61Å) and C II(1037.02Å) emission feature. JKH do not present
the parameters used for this scaling, shifting, and combination.

From first principles, the intrinsic relative strength of the C II 1335.7Å
line to the C II 1037.02Å line is hard to predict. The ratio of $gf$$\lambda$
values gives a factor of 5.3 for the $\lambda\lambda$1335.7/1037.02 ratio, and
the Boltzmann factor ratio could be 2.4 at a temperature of 40000K suggesting
the 1037.02Å line could be 13 times weaker than the 1335.71Å line under
optically thin conditions in LTE, and not 1.5 times weaker as JKH propose.
However unknown optical depths and non-LTE conditions occur in the formation
of resonance lines of C II causing them to appear weaker and decrease the
$\lambda\lambda$1335.7/1037.02 ratio. These issues confound any meaningful
scaling. Moreover, the line profiles themselves change with time as
demonstrated earlier. Comparison of HST spectra from 2000 to FUSE spectra from
2003 (cf. Figure 13) clearly shows the substantial changes in the emission
profiles of ultraviolet lines on the ‘red’ sides ($+$100 to $+$400 km s-1)
between these three years. It is obvious from HST spectra of C II in 2000,
2002, and 2010 (Figure 18) that the critical region of the 1335.7Å line
profile used to assess absorption, namely the region from $+$35 to $+$85 km
s-1, varies in both absolute flux as well as slope. The claim made by JKH —
that this scaled, shifted, blended feature does not exhibit strong absorption
on the red side caused by the hot O VI wind — is unjustified, in particular
given the substantial changes in both the emission and the wind.

(4) JKH suggest that two H2 lines arising from the circumstellar disk exhibit
behavior that suggests a cool wind is present, but not a hot wind. Their
interpretation does not agree with the geometry of TW Hya and its disk, nor
with the known characteristics of a wind from a dwarf star.

The H2 line at 1333.85Å occurs $-$165 km s-1 from the C II line at 1334.5Å (at
$-$430 km s-1 in Figure 18). JKH claim that this H2 feature is absorbed by the
C II wind because their model overpredicts its strength in their spectrum
taken in 2000. By contrast, an H2 line lying at a velocity of $-$165 km s-1
from the C IV 1548Å feature does not suffer wind absorption when compared to
their model (Herczeg et al. 2002). Thus JKH conclude that a cool wind exists,
but not a hot wind.

Many spectra subsequent to the one used by JKH reveal conflicting patterns
(see Fig. 18). In 2002, the C II wind extends to $-$251 km s-1, which would
weaken not only the line located at 1333.85Å, but should weaken another H2
line located at $-$225 km s-1 (corresponding to $-$500 km s-1 in Figure 17),
but it does not. Additionally, in the presence of a strong C II wind, another
H2 line emerges at 1335.2Å ($\sim-$125 km s-1 in Fig. 18; the R(2) 0-4
transition) which did not appear previously. All of these spectra indicate
optically thick winds in C II are present and the wind opacity is variable.
Furthermore, the H2 lines behave in apparently unpredictable fashion – perhaps
because they are produced through photoexcitation by variable Lyman-$\alpha$
emission. With a more precise location of the H2 emission, the behavior of the
molecular lines might aid in the wind description. However, since these lines
are pumped by H-Ly$\alpha$, which itself varies, trying to pin down a
consistent model may prove challenging. Moreover, serious considerations
discussed earlier in (3), rule out the putative effects of wind absorption.

Trying to assess potential effects of a wind upon the observed H2 emission
involves many factors. First, the line opacity in the wind must coincide in
wavelength with the H2 emission from the disk. H2 is believed to be produced
less than 2 AU from the star (Herczeg et al. 2002) where it originates in the
gaseous component of the circumstellar disk around TW Hya. Recent ALMA
observations suggest the dust-depleted cavity lies within 4 AU of the star
(Rosenfeld et al. 2012). Ions in a radially-moving stellar wind passing over a
distant disk will absorb at lower velocities than when observed against the
nearly pole-on star itself. Rosenfeld et al. (2012) suggest the TW Hya disk
might have a warp of $\sim$8 degrees, such that the effective velocity of any
absorption would occur at a still lower velocity. Thus the wind may not have
significant opacity at the position of the H2 emission when passing across the
disk at a low angle. Moreover, any channeling, flux-tube expansion, or
inhomogeneities (all of which occur in the solar wind) will modify the amount
of absorption and the velocity at which it occurs. A more fundamental concern
appears to be the ionization stage in the wind itself. If the stellar wind at
2 AU does not contain C II or C IV, then absorption will not occur. From very
detailed knowledge about the wind of one star, our Sun, the low species ions
are effectively absent at a distance of 1 AU (Landi et al. 2012) and the wind
contains the ionization stages of the corona. Thus at a distance of 1 AU or
larger, the wind of TW Hya, if similar to the solar wind, would appear to lack
the ions that JKH invoke for their arguments.

The presence of a hot wind from TW Hya is not unexpected. The Sun is a low
mass star, with weak magnetic fields and it possesses an outer atmosphere with
temperatures of a few MK. Wave-heating mechanisms (Kohl et al. 2006) or
possibly nano-flares resulting from reconnection (Tripathi et al. 2010) can
heat the corona. Self-consistent models for hot winds have been constructed
for accreting T Tauri stars (Cranmer 2008). These winds are turbulence driven,
and the turbulence has two sources: interior convection and ‘ripples’ from
nearby accretion shocks. Both observations and theory offer much evidence that
low-mass cool stars can support a hot outer atmosphere.

We conclude that the stellar wind of TW Hya is a hot wind, evidenced by the
ultraviolet spectra as originally suggested (Dupree et al. 2005a).

##  9 Discussion and Conclusions

The many spectra of TW Hya reported here document the source and components of
emission line features and the variability of the accretion process,
subsequent atmospheric heating, and wind expansion. We find:

1\. Time-domain spectroscopy reveals symmetric H$\alpha$ profiles in
approximately half of our spectra and indications of outflowing wind
asymmetries in the remaining spectra. This spectroscopic example indicates
that the source of the broad emission feature arises from the TW Hya star
itself, and can be identified with the post-shock cooling zone. Broad lines
ranging from X-ray through infrared transitions can be straightforwardly
explained by the presence of a variable turbulent post-shock cooling volume
producing the emission which can be modified by wind scattering and/or
material infall. Earlier evidence of a logical progression in line shape and
flux following an X-Ray accretion event also supports this interpretation
(Dupree et al. 2012). Models are now needed for the temperature structure and
emergent emission from this post-shock region. A recent MHD simulation
(Matsakos et al. 2013) of a two-dimensional accretion shock impacting a
chromosphere suggests that substantial chaotic motions can result in the low
chromosphere. These need to be followed up with calculation of the radiation
emitted from the cooling plasma.

2\. Absorption features at negative velocities in the H$\alpha$ line noted by
Alencar et al (2002), are demonstrated with the frequent continuous spectral
coverage reported in this paper, to be stable in velocity during a night’s
observation. Moreover, their stable behavior and opacity characteristics
suggest they appear to arise from the silhouette of an accretion stream high
in the stellar chromosphere. At positive velocities, absorption in H$\alpha$
mimics the infall speeds detected in He I.

3\. Current models (Muzerolle et al. 2001; Kurosawa & Romanova 2012, 2013)
that consider the optical emission from accretion ‘funnels’ do not agree with
the observed line profiles in TW Hya. The observed profiles lack the shift,
the asymmetries and/or absorption predicted from accretion stream models. In
addition, many models do not include a stellar wind, which is clearly
indicated in the observed line profiles, nor emission from a post-shock
cooling region.

4\. A hot wind, and its variable speed and opacity are revealed by the UV and
far-UV resonance lines with spectra spanning a decade in time. Previous
arguments offered (JKH) against the presence of a hot wind are shown to be
unpersuasive. Broad line profiles and variability support both formation in a
turbulent medium and the presence of a hot wind. This hot wind is a natural
extension of the wind indicated by the near-IR helium line profile.

5\. The He I near-IR transition at 10830Å shows less night-to-night absorption
variability than H$\alpha$, suggesting a large-scale wind outflow from the
star. The emission in the helium line correlates inversely with the velocity
of the accreting inflowing material consistent with a simple model of dipole
accretion, and formation of the emission in the turbulent post-shock cooling
zone. The stellar wind velocity appears related to the strength of helium
emission providing direct spectroscopic evidence that the accretion process
leads to wind acceleration.

We thank the anonymous referee for his/her comments that improved the
manuscript. The authors gratefully acknowledge the helpful support from
Gemini-S astronomers, the staff at Magellan, and KECK II while acquiring these
spectra. Observers at FLWO were a great help on short notice. This research
has made use of NASA’s Astrophysics Data System Bibliographic Services. Some
of the data presented here was obtained from the Mikulski Archive for Space
Telescopes (MAST). STScI is operated by the Association of Universities for
Research in Astronomy, Inc. under NASA contract NAS5-26555. Support for MAST
for non-HST data is provided by the NASA Office of Space Science via grant
NNX09AF08G and by other grants and contracts. We wish to extend special thanks
to those of Hawaiian ancestry from whose sacred mountain of Mauna Kea we are
privileged to conduct observations. Without their generous hospitality, the
Keck results presented in this paper would not have been possible.
Facilities: Gemini:South (PHOENIX), Magellan:Clay (MIKE), HST(STIS), FUSE,
FLWO:1.5m (TRES), Keck II (NIRSPEC)

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Figure 1: Single H$\alpha$ spectra of TW Hya selected from ten years of
measurements from 2004 to 2013. Note that there are two different scales on
the y-axis. Observations in 2005 Jul, 2006 Apr, and 2011 Apr indicate stronger
H$\alpha$. The narrow absorption feature near $+$50 km s-1 is due to water
vapor. All of these spectra suggest that H$\alpha$ is centered on the stellar
radial velocity, and the profile is substantially modified by wind absorption
causing the apparent ‘red’ enhancement. Figure 2: The observed H$\alpha$
spectrum from 2010 Jun 30 (solid line) with two models overlaid from
Muzzerolle et al. (2001) that assume either no line-damping or an arbitrary
damping parameter. These models assume formation of H$\alpha$ in an accretion
stream which can be seen by the weaker positive velocity side of the profile
as compared with the negative side, and slight subcontinuum absorption when
damping is absent. A stellar wind is not included in the models accounting in
part for the excess emission of the model on the negative velocity side. Also
the model profiles appear ‘pointed’ which is not found in the observed
profiles (see also Figure 1). Figure 3: Magellan/MIKE spectra over 4 nights in
2004 showing the appearance of wind structures stable in velocity. The time
interval between the first and last spectrum is shown. Each spectrum has been
offset for display. The broken vertical lines mark the constant velocity of
the absorption feature during each night. The notch in the H$\alpha$ emission
at $+$50 km s-1 results from water vapor absorption. Figure 4: Magellan/MIKE
spectra of the H$\alpha$ line over 3 consecutive nights in April 2006. The
narrow absorption near $+$50 km s-1 is caused by water vapor. More spectra
were obtained than displayed here and a gray scale representation of all
spectra is shown in Figure 5.

Figure 5: Left panel: Gray scale representation of the H$\alpha$ profiles of
TW Hya over 3 consecutive nights in April 2006 where each spectrum has been
divided by the maximum flux at each wavelength in order to display absorption
features. During Night 2, absorption at $-$50 km s-1 is replaced by a
discontinuous jump to a new absorption feature at $-$200 km s-1 (at $\sim$29
hrs). Right panel: Profiles during night 2 of the gray scale representation
above. Note the weakening of the absorption at $-$50 km s-1 during the night,
and the abrupt appearance of another absorption feature at $-$200 km s-1 in
the middle of the night which becomes prominent at the end of the night. Broad
absorption extending from $\sim+$150 to $\sim+$400 km s-1 arising from
infalling material is variable, and the velocities are commensurate with the
velocities measured in He I. Figure 6: Balmer series over 4 consecutive nights
in 2007 (Feb. 26 – Mar. 1). The y-axis represents a continuum normalized flux.
Profiles from night 4 (2007 March 1) have the negative velocity segments
overlaid on the positive velocity side and marked by a broken (red) line. The
effects of wind absorption can be seen when the broken line lies below the
solid line; effects of absorption by the accretion stream can be seen at
$\sim+$200 km s-1 where the broken line lies slightly above the solid line.
Figure 7: Typical H$\delta$ line profile ($\lambda$4101.73) over 4 consecutive
nights in 2008. The broken line represents the negative velocity wing of the
first night’s observation that has been reflected around the axis at zero
velocity. This illustrates the slight broad absorption present at velocities
+50 to +250 km s-1. Profiles from the following 3 nights are symmetric.
Subcontinuum absorption predicted by magnetospheric accretion models is
missing in the positive velocity wing of the line. Figure 8: Spectra of He I
10830Å spanning 8 years. Substantial changes in the emission flux level, the
wind opacity, the wind speed, and the inflowing absorption are evident.

Figure 9: KECK II NIRSPEC spectra (2005) and Gemini-S PHOENIX spectra (2009,
2010) of TW Hya on successive nights. The outflowing wind remains relatively
stable as compared to both the emission and the inflowing material.

Figure 10: Left panel: The magnetic dipole accretion model predicts the
relationship between the accretion rate and the terminal free-fall velocity. A
slower infall velocity results from a smaller radius of the inner
circumstellar disk, and a larger filling factor in the stellar atmosphere
produces a higher mass accretion rate. Parameter ranges constrained by CHANDRA
Ne IX diagnostics are shown here. The $\times$ symbols denote the average
CHANDRA spectrum, while the other symbols mark 3 individual CHANDRA pointings
(cf. Brickhouse et al. 2012). Right panel: Equivalent width of the emission
component of the He I line ($\lambda$10830) as a function of the terminal
velocity indicated by the inflowing subcontinuum absorption. These
observations display general agreement with the magnetic dipole accretion
model. The emission in the post-shock cooling zone increases because of
increased accretion at lower terminal free-fall velocities. Figure 11: The
terminal outflow velocity as a function of the emission component of the He I.
Except for one outlier (2007 Mar. 1), fast outflow is correlated with
increased emission. This is not inconsistent with the conclusion that the
turbulent post-shock region can affect the wind speed.

Figure 12: UV and far-UV spectra of major resonance lines. Peak fluxes are
scaled to match to illustrate the shape of the profiles. Left panel: HST:STIS
spectra of N V ($\lambda$1239) and C IV($\lambda$1548) taken in the year
2000\. Right panel: FUSE spectra of C III ($\lambda$977) and O VI taken three
years later in 2003 (Dupree et al. 2005a). Note that the profiles observed at
the same time are similar one to another with the exception of different
absorption levels on the negative velocity side (see text for explanation).
Figure 13: Comparison of resonance lines at different times showing the
dramatic change in emission at positive velocities. The second member of the C
IV doublet appears at $+$500 km s-1. See text for discussion of these
profiles. Peak fluxes are scaled in this figure to illustrate the shape of the
line profiles.

Figure 14: C II (1335Å) and C IV (1548Å) multiplets showing the changes in
both intrinsic line strength and wind opacity, even on short time scales, such
as the 6 days separating 2010 Jan 29 and 2010 Feb 4. The C II line is
generally narrower on the positive velocity side than the C IV which is
similar to the high temperature N V profile. The broken vertical lines mark
the rest velocities of the doublets.

Figure 15: FUSE time-tagged spectra of C III (977Å) and O VI (1031.91Å) taken
through the LWRS aperture from the SiC2A and LiF1A detector measured over a
32.2 hour span 2003 Feb 20-21. The variability at positive velocities is more
substantial in both lines than at negative velocities. The opacity of the hot
wind as measured in the O VI line remains relatively constant. Flux
measurements from the individual spectra are shown in Figure 16.  Figure 16:
Fluxes measured from the FUSE line profiles shown in the previous figure
display similar behavior in the C III and O VI emission during several hours
of the second pointing. During the second pointing, the absolute flux values
are reversed from the first pointing, and the line fluxes are correlated.

Figure 17: The solid lines show the ‘CoolCAT’ optimum reduction (upper panel) of the same STIS spectrum as that used by JKH (lower panel, O59D01030). The broken lines (red in online version) represent the profile of the 1548Å multiplet multiplied by 0.5 and shifted by the velocity equivalent of the wavelength separation of the doublet. Wind absorption by the 1550Å component is indicated by the weaker extended positive velocity wing of the 1548Å transition (broken line). Wind absorption occurs systematically between $+$175 and $+$375 km s-1 of the scaled 1548Å line in the velocity scale of the figure. Additionally the effect of increased opacity in the 1548Å line is evident in the region $-$125 to $+$25 km s-1, where the broken line lies systematically below the solid line. Thus, two signatures of wind absorption can be noted in these doublet profiles: absorption of the red wing of 1548Å by the blue wing of 1550Å, and the increased scattering of the 1548Å line in the wind as compared with 1550Å. Figure 18: HST/STIS spectra of TW Hya in the region of the C II 1335Å doublet in 2000, 2002, and 2010 showing the variabilities in C II emission, H2 emission and wind speed and opacity. The short wavelength component ($-$275 km s-1) of the doublet is bifurcated by interstellar absorption. The broken vertical lines mark the region ($+$35 to $+$85 km s-1) where JKH postulated the lack of absorption. This section of the profile is obviously substantially variable in both flux and slope of the emission line. The arrows indicate the extent of wind absorption for one member of the doublet. Note that the $H_{2}$ lines vary, but not in any systematic way with wind speed or opacity. In 2002, when the wind opacity is substantial between $-$100 and $-$250 km s-1, another $H_{2}$ emission line appears (R(2)0-4 at 1335.2Å at $\sim-$125 km s-1).  Table 1: TW Hya Observations Spectral | Instrument | Date | Data file/ | Exp. time
---|---|---|---|---
Region |  |  | spectra | (s)
Far UV | FUSE | 2003 Feb 20 | C0670101aaPI: J. Linsky | 15221
| FUSE | 2003 Feb 21 | C0670102aaPI: J. Linsky | 15465
UV | HST/STIS/E140M | 2000 May 7 | O59D01030bbCoolCAT (casa.colorado.edu/ $\tilde{}$ayres/CoolCAT) was derived from this spectrum. | 2300
| HST/STIS/E140H | 2002 July 19 | StarCATccStarCAT can be found at archive.stsci.edu/prepds/starcat (Ayres 2010). | 4355
| HST/STIS/E140M | 2010 Jan 29 | OB3R07060ddPI: N. Calvet | 3058
| HST/STIS/E140M | 2010 Feb 04 | OB3R08060ddPI: N. Calvet | 3058
| HST/STIS/E140M | 2010 May 28 | OB3R09060ddPI: N. Calvet | 3058
Optical | Magellan/CLAY/MIKE | 2004 Apr 27–30 | 19 | 60-300
| Magellan/CLAY/MIKE | 2005 Jul 24–27 | 4 | 60-300
| Magellan/CLAY/MIKE | 2006 Apr 15–17 | 19 | 60-300
| Magellan/CLAY/MIKE | 2007 Feb 26–Mar 1 |  $\sim$360 | 45-360
| Magellan/CLAY/MIKE | 2007 Jul 23–25 | 4 | 45-60
| Magellan/CLAY/MIKE | 2008 May 25–26 | 7 | 45-240
| Magellan/CLAY/MIKE | 2008 Jul 13–16 | 9 | 35-180
| Magellan/CLAY/MIKE | 2009 May 4–5 | 5 | 45-180
| FLWO/TRES | 2009 Jun 10–12 | 3 | 720
| Magellan/CLAY/MIKE | 2010 Jun 30–Jul 2 | 3 | 45
| FLWO/TRES | 2011 Apr 11–15 | 5 | 1200
| Magellan/CLAY/MIKE | 2013 Jul 19–25 | 5 | 30-180
Near IR | KECK II/NIRSPEC | 2002 May 20 | 2 | 200
| KECK II/NIRSPEC | 2005 May 17–18 | 4 | 200
| Gemini-S/PHOENIX | 2007 Mar 1 | 1 | 900
| Gemini-S/PHOENIX | 2009 Jun 10-12 | 3 | 300
| Gemini-S/PHOENIX | 2010 Mar 25–27 | 3 | 300