# Fast Winds and Mass Loss from Metal-Poor Field Giants111Data 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

A. K. Dupree  Harvard-Smithsonian Center for Astrophysics, 60 Garden Street,
Cambridge, MA 02138<EMAIL_ADDRESS>Graeme H. Smith  University of
California Observatories/Lick Observatory, University of California, Santa
Cruz, CA 95064<EMAIL_ADDRESS>Jay Strader222Hubble Fellow Harvard-
Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138
<EMAIL_ADDRESS>
###### Abstract

Echelle spectra of the infrared He I $\lambda$10830 line were obtained with
NIRSPEC on the Keck 2 telescope for 41 metal-deficient field giant stars
including those on the red giant branch (RGB), asymptotic giant branch (AGB),
and red horizontal branch (RHB). The presence of this He I line is ubiquitous
in stars with $T_{eff}\gtrsim$ 4500K and $M_{V}$ fainter than $-$1.5, and
reveals the dynamics of the atmosphere. The line strength increases with
effective temperature for $T_{eff}\gtrsim$ 5300K in RHB stars. In AGB and RGB
stars, the line strength increases with luminosity. Fast outflows ($\gtrsim$
60 km s-1) are detected from the majority of the stars and about 40 percent of
the outflows have sufficient speed as to allow escape of material from the
star as well as from a globular cluster. Outflow speeds and line strengths do
not depend on metallicity for our sample ([Fe/H]= $-$0.7 to $-$3.0) suggesting
the driving mechanism for these winds derives from magnetic and/or
hydrodynamic processes. Gas outflows are present in every luminous giant, but
are not detected in all stars of lower luminosity indicating possible
variability. Mass loss rates ranging from $\sim 3\times 10^{-10}$ to $\sim
6\times 10^{-8}$ $M_{\sun}\ yr^{-1}$ estimated from the Sobolev approximation
for line formation represent values with evolutionary significance for red
giants and red horizontal branch stars. We estimate that 0.2 M☉ will be lost
on the red giant branch, and the torque of this wind can account for
observations of slowly rotating RHB stars in the field. About 0.1–0.2 M☉ will
be lost on the red horizontal branch itself. This first empirical
determination of mass loss on the RHB may contribute to the appearance of
extended horizontal branches in globular clusters. The spectra appear to
resolve the problem of missing intracluster material in globular clusters.
Opportunities exist for ’wind smothering’ of dwarf stars by winds from the
evolved population, possibly leading to surface pollution in regions of high
stellar density.

stars: chromospheres — stars: Population II — stars: winds, outflows

##  1 Introduction

The assumption of mass loss from stars evolving on the red giant branch of
globular clusters has yet to be tested through direct detection of winds. This
assumption remains one of the major concerns in the evolution of low mass
stars and may be related to the second-parameter problem (Sandage & Wildey
1967): differing horizontal branch morphology between globular clusters of the
same metallicity and age. Various explanations have been offered for
differences in the horizontal branch morphology including intrinsic
dispersions in the amount of stellar mass loss, rotation, or deep mixing,
environmental effects possibly correlated with cluster mass or central
density, heterogeneities in He abundance possibly as a result of cluster
pollution by intermediate-mass asymptotic giant branch (AGB) stars, or of the
infall of planets onto cluster stars (Buonanno et al. 1993; Buonanno et al.
1998; Catelan et al. 2001; Sills & Pinsonneault 2000; Recio-Blanco et al.
2006; Sandquist & Martel 2007; Soker et al. 2001b; Sneden et al. 2004;
Sweigart 1997; Peterson et al. 1995; Ventura & D’Antona 2005). Most recently
with the advent of infrared photometry discussed below, and the identification
of multiple populations on the main sequence of globular clusters with the
Hubble Space Telescope (Anderson 2002; Bedin et al. 2004; Piotto et al. 2007),
renewed attention is focussing on the mass loss process.

A related issue is the puzzling absence of the material lost from the red
giants in a globular cluster (Tayler & Wood 1975). This gas is expected to
accumulate in clusters between sweeps through the galactic plane, but
detection has proved elusive. Unconfirmed measures of 21-cm H-line emission
exist for NGC 2808 (Faulkner et al. 1991). The radio dispersion of milli-
second pulsars located in 47 Tuc and M15 hinted at an enhanced electron
density in the intracluster medium (Freire et al. 2001). Deep searches for H I
in 5 clusters gave one firm detection of H I emission, and possibly 2 others,
but the amount of mass inferred from this is a few orders of magnitude less
than expected (van Loon et al. 2006). A search for intracluster dust in 12
globular clusters with the Far-Infrared Surveyor on AKARI failed to detect
emission except possibly in one cluster (Matsunaga et al. 2008). There is a
notable lack of detections of intracluster material with the Spitzer Space
Telescope. Only 2 clusters of the many observed with Spitzer have ’possible’
detections of intracluster material (Boyer et al. 2006; Barmby et al. 2009).
One suggestion for removal of intracluster material is ram-pressure stripping
by the galactic halo (Frank & Gisler 1976, Okada et al. 2007), but Spitzer
observations can not confirm the expected relation between cluster kinematics
and the presence (or upper limit) of dust (Barmby et al. 2009).

Results from mid-infrared observations of metal-deficient field giants and
globular clusters with IRAS, the ISO satellite, the Spitzer Space Telescope,
and AKARI (Smith 1998; Origlia et al. 2002, 2007; Evans et al. 2003; Boyer et
al. 2006; Ita et al. 2007; Boyer et al. 2008) suggest that some giants have
produced circumstellar dust that could result from stellar winds. Since not
all giants display excess infrared emission, mass loss associated with dust
appears to be episodic (Origlia et al. 2007; Mészáros et al. 2008). Caloi and
D’Antona (2008) further suggest that the mass loss rate might be ‘sharply’
peaked at one value along the red giant branch. In the metal-rich open cluster
NGC 6791, the presence of low mass white dwarfs led Kalirai et al. (2007) to
conclude that mass loss is enhanced in high metallicity environments, although
that suggestion appears to be controversial (van Loon et al. 2008; Bedin et
al. 2008). Concurrently with the above results, the possibility of multiple
episodes of star-formation, and self-pollution in globular clusters is
receiving increased attention as a means to explain chemical variations and
multiple branches in the color-magnitude diagram of clusters (Lee et al. 1999;
Pancino et al. 2000; Anderson 2002; D’Antona et al. 2002; Piotto et al. 2007;
Kayser et al. 2008). This suggestion frequently resorts to AGB stars with mass
$>~{}$3M☉ ‘polluting’ either the surface of the cluster stars or the
environment in which a second generation of stars subsequently forms. The wind
velocities from the AGB stars must be slow so that material will not escape
the cluster. A long-standing suggestion envisions that pulsation near the top
of the AGB may degenerate into relaxation oscillations, during which mass loss
or envelope ejection occurs in rapid fashion, dubbed a ‘superwind’ (Renzini
1981; Bowen and Willson 1991). Another conjecture (Soker et al. 2001a) calls
for a superwind during the immediate post-RGB phase to explain gaps on the
horizontal branch. Clearly, both observations and theory currently allow a
great variety in both the presence and character of mass loss from cluster
stars. The spectra of He I $\lambda$10830 presented here can address these
questions.

##  2 Spectroscopic Diagnostics of Winds

Spectra of stars contain many features that can be used to detect winds
directly, but their diagnostic properties are related to the specific
conditions in a stellar atmosphere. In a cool giant star, atoms and ions form
at different levels in the chromosphere, loosely tied to the local electron
temperature which is increasing with height. Additionally, for strong lines
such as H$\alpha$, Ca II H and K, and Mg II, the profile is formed over an
extended range of atmospheric layers which results from differing opacities
across the line itself. Calculations for metal-poor giants demonstrate (Dupree
et al. 1992a; Mauas et al. 2006) the locations in the atmosphere where
commonly observed lines are formed. Mass flows can produce asymmetric profiles
of H$\alpha$ and Na D, as well as emission asymmetries and velocity shifts in
the reversed absorption core of Ca II and Mg II. These features have been
measured in red giants in many globular clusters (Cacciari et al. 2004;
McDonald and van Loon 2007; Mészáros et al. 2008, 2009b) and in metal-poor
field stars (Smith & Dupree 1988; Dupree & Smith 1995; Dupree et al. 2007).
However the velocities inferred from the optical profiles are generally less
than the escape velocity from the star, and can not truly be identified as
stellar winds.

We note the difference between the use of ’outflow’ and ’wind’ in this
context. By ’outflow’, we mean the presence of a velocity field within the
region of a chromosphere sampled by our spectral diagnostic in which a flow of
material occurs that is moving away from the stellar photosphere. The term
’wind’ is reserved for the particular case of an outflow in which the outflow
velocity exceeds the escape velocity within the region of the chromosphere
sampled by our spectral diagnostic.

A good diagnostic of winds is the near-infrared He I 10830Å line ($1s2s\
^{3}S\ -\ 1s2p\ ^{3}P$) which models show (Dupree et al. 1992a) is formed
higher in the metal-poor atmosphere than H$\alpha$ and Ca II K. Thus it might
be expected to trace out higher velocities, where the outflow becomes a wind,
than the optical diagnostics. Additionally, the lower level of this transition
is metastable, and is not closely linked to local physical conditions in the
wind, so it can absorb photospheric radiation and map out an expanding wind in
a luminous star.

The lower level of the $\lambda$10830 multiplet ($1s2s\ ^{3}S_{1}$) lies 19.7
eV above the ground state of He I and can be populated directly by collisions
from the ground state ($1s^{2}\ {}^{1}S$) although the rate is far less than
for an allowed transition. High temperatures ($\gtrsim$ 20,000 K) are
generally required. The ${}^{3}S$ level can also be populated by recombination
from the continuum. This latter pathway has long been studied especially in
the Sun and other cool stars (Goldberg 1939; Harvey & Sheeley 1979; Zirin
1975, 1982) because a source of EUV or X-ray radiation can photoionize He I
from the ground state or the ${}^{3}S$ level and the helium ion preferentially
recombines to the triplet state followed by subsequent cascade to the lower
level of the $\lambda$10830 transition. Thus stars with strong X-ray emission
display enhanced helium absorption in $\lambda$10830 (O’Brien & Lambert 1986;
Zarro & Zirin 1986; Sanz-Forcada & Dupree 2008). Depopulation of the ${}^{3}S$
state occurs at high densities with collisions to the ${}^{1}P$ level. Because
${}^{3}S$ has a long lifetime for decay to the ground state, (i.e. this level
is metastable), a significant population can build up, and provide an
opportunity for scattering of near-IR photons from the line itself or
continuum photons from the photosphere. If the chromosphere is expanding, this
transition can trace out the wind velocity as it scatters radiation while
being carried along in the expansion. An additional advantage of this
transition is that the profile can not be compromised by interstellar or
circumstellar absorption since it is not a resonance line.

The first detection of a wind in a metal-deficient star using the
$\lambda$10830 line was made (Dupree et al. 1992a) in the bright field giant,
HD 6833 where an outflow of 90 km s-1 was discovered $-$ a value comparable to
the chromospheric escape velocity. Subsequently, Smith et al. (2004)
identified He I absorption from one warm AGB star in the globular cluster M13
in addition to two other metal-poor field giants. The short wavelength
extension of the lines in these stars reached 90 to 140 km s-1 – again fast
enough to escape a chromosphere and also a globular cluster. A stellar
$T_{eff}$ greater than 4600K appeared required to populate the lower level of
the He I atom; thus Smith et al. (2004) suggested that the coolest red giants
can not produce this transition. Indeed, $\lambda$10830 was not detected in
the 5 coolest red giants observed in M13. While globular cluster stars
themselves remain ideal targets, the metal-deficient field giants are
brighter, more accessible to current instrumentation, and can act as
surrogates for cluster stars. We report here on the high-resolution
spectroscopy of the He I $\lambda$10830 line in 41 such field stars.

##  3 Observations and Reductions

The objective in this investigation was to study the systematics of the He I
line among evolved Population II stars in a variety of evolutionary states.
This goal suggests an observational program concentrating on halo field stars
rather than globular cluster stars, since in the latter only the upper regions
of the red giant and asymptotic giant branches can be studied at high signal-
to-noise, even with the NIRSPEC instrument on the Keck 2 telescope. Spectra of
the He I line for several red giants in the cluster M13 were published by
Smith et al. (2004). In our study, halo field stars were chosen from the lists
of Bond (1980) and Beers et al. (2000) with an effort to achieve a good
sampling in the red giant branch, red horizontal branch, and asymptotic giant
branch phases of evolution. No selection was made on the basis of metallicity
or proper motion, although certain radial velocities were avoided in order to
prevent overlap of the He I line with telluric absorption features. Some
chemically peculiar stars in the form of CH stars were included in the sample.
To facilitate high signal-to-noise spectroscopy, only stars with apparent
magnitudes of $V<11$ were observed, and most have $V<10$. Thus this sample
comprises relatively nearby halo stars, although most are beyond the limit at
which the Hipparcos satellite provided reliable parallaxes. In order to study
evolved stars, most of our targets have absolute magnitudes of $M_{V}<+2$, and
no subdwarfs were included in the program.

The spectra of 41 metal-poor halo field giants (Table 1) were obtained during
1.5 nights of observation in May 2005 using NIRSPEC (McLean et al. 1998, 2000)
on the Keck 2 telescope. Observations were made using the echelle cross-
dispersed mode of NIRSPEC with the NIRSPEC-1 order-sorting filter and a slit
of $0.43^{\arcsec}\times 12^{\arcsec}$ giving a nominal resolving power of
23,600. The long-wavelength blocking filter was not used in order to minimize
unwanted fringing. Total integration times accumulated for each star are
listed in Table 2; these times are generally broken into two shorter exposures
in the NOD-2 positions. Calibration exposures consisted of internal flat-field
lamps, NeArKr arcs, and dark frames. Spectra of rapidly rotating hot stars
obtained at airmasses similar to the target objects were used to identify, and
minimize or eliminate night sky emission lines.

Data reduction was performed using the REDSPEC package (McLean et al. 2003)
which was written specifically for NIRSPEC. Following dark subtraction and
flat fielding, the data frames were spatially rectified. Wavelength
calibration was performed using NeArKr arc lamp spectra taken after each
science observation. For this study, we extracted only order 70 (wavelength
coverage $\sim 1.079-1.095$ $\mu$m), which contains the He I line at 1.0830
$\mu$m. The strong Si I photospheric absorption line at 1.0827$\mu$m was used
to estimate the stellar radial velocity with an uncertainty of $\sim$5 km s-1.
Values of the radial velocity (RV) are given in Table 2.

The spectrum of each giant was normalized to a continuum determined by fitting
a 5th order cubic spline to the wavelength range 1.082 - 1.092 $\mu$m. The
wavelength scales of the spectra were then shifted onto the photospheric rest
frame of each star by applying a zero-point wavelength shift determined from
measuring the wavelength of the nearby photospheric 10827.09Å Si I line. The
equivalent width of the helium line was measured from the continuum normalized
spectra using the IRAF333Image Reduction and Analysis Facility (IRAF) written
and supported by the IRAF programming group at NOAO, operated by AURA under a
cooperative agreement with the NSF. (http://iraf.noao.edu/iraf/web) task
‘splot’ and measuring the width directly, or, if blended with the adjacent Si
I line, deconvolving the blend with Voigt profiles and dividing the total
equivalent width measured directly into appropriate fractions obtained from
the deconvolution. The photometric colors, extinction, and evolutionary state
as inferred for the target stars are included in Table 1. The footnotes to
this Table contain the references to the tabulated quantities. The assignment
of the evolutionary state of our targets is based on the Strömgren $c_{1}$ vs.
$b-y$ diagram and the color magnitude diagram ($M_{V}$ vs $B-V$). Parameters
of the He I line are given in Table 2. About half (21) of the 41 targets are
located on the red giant branch (RGB), and one (HD135148) is identified as a
CH star. Red horizontal branch (RHB) stars comprised 11 objects, and
asymptotic branch stars (AGB) made up 6 targets. Two subgiants and one semi-
regular red variable (TY Vir) completed the sample. The spectra of these stars
are shown in Figure 1 (RGB stars), Figure 2 (RHB stars), Figure 3 (AGB stars),
and Figure 4 (subgiant stars).

The stars in our sample show three basic types of He I line behavior: (i) a
helium line with a pure absorption profile, (ii) a P Cygni type profile in
which an absorption profile is paired with an emission feature to longer
wavelengths, and (iii) no helium line at all - neither in absorption or
emission. Inspection shows that the He I line is generally broader than the
neighboring Si I photospheric absorption line which is expected since He I
arises in the higher temperature chromosphere. In many cases where a He I
absorption profile is present, this profile extends to shorter wavelengths
providing evidence of a chromospheric outflow (see, for example, the spectra
of stars BD $+$30°2611 and HD 122956 in Figure 1, BD $-$03°5215, HD 119516 in
Figure 2, and HD 121135 and HD 107752 in Figure 3).

The star HD 135148, classified as a CH star, and identified as a spectroscopic
binary (Carney et al. 2003) exhibits (Figure 5) a substantial P Cygni profile
with deep absorption almost to zero flux at velocities $\sim-$60 km s-1, and
extending to $-$115 km s-1. The spectrum of the coolest target of our sample,
HD104207 (GK Com) is shown in Figure 6 where many photospheric absorption
features of neutral atoms appear in addition to a weak P Cygni feature of He
I. The value of [Fe/H]=$-$1.93 in this star is similar to many other stars in
our sample without such an array of neutral species in their spectra, thus
dramatically illustrating the effects of low effective temperature.

##  4 Discussion

The targeted sample includes $(B-V)$ colors ranging from 0.6 to 1.6, and spans
6 magnitudes, reaching from the tip of the RGB to several subgiant stars. Most
of the stars showed a He I $\lambda$10830 feature the presence of which is
shown as a function of position in the $M_{V}$ versus $(B-V)_{0}$ and Teff
diagrams in Figure 7\. The targets from M13 and the metal-deficient field
giants reported earlier (Dupree et al. 1992a; Smith et al. 2004) have been
added to this figure and their parameters are given in Table 3. Many of the
more luminous stars at $M_{V}=0$ and brighter, exhibit emission resulting from
scattering of the line above the stellar limb which is normal in large stars
with extended chromospheres. Since the He I line is formed at chromospheric
temperatures ($\sim$10,000$-$20,000K), it might be expected to vanish in the
coolest objects where the chromosphere does not attain sufficiently high
temperatures. A region exists in the color magnitude diagram, with magnitude
brighter than $-$1.5 and with $(B-V)_{0}>1.1$ where the $\lambda$10830 line
does not generally appear either in emission or absorption. Our earlier search
for He I along the red giant branch in M13 revealed absorption in an AGB star,
IV-15, near $(B-V)_{0}$ =1.02, but not in any of 5 cooler stars on the red
giant branch (Smith et al. 2004). The presence of a $\lambda$10830 line among
Population II giants as a function of position in a color-magnitude diagram
differs from Population I stars (O’Brien & Lambert 1986; Lambert 1987), where
the helium emission disappears near spectral type M1 in giants and supergiant
stars, corresponding to $T_{eff}$ of 3780 K (Tokunaga 2000). Based on our
earlier sample (Smith et al. 2004), we noted that Population II giants with
$T_{eff}$ less than 4600K did not show helium. However the larger sample
presented here contains several stars with $T_{eff}$ less than 4600K, and they
exhibit the $\lambda$10830 line. Two of these objects are somewhat anomalous
red giants. HD 104207 (GK Com) is the coolest star in the sample, and a semi-
regular variable. It is plausible that the atmosphere cycles through heating
and cooling phases, producing the helium line at certain times. The other
star, HD 135138, is a CH object, a spectroscopic binary with a degenerate
secondary star both of which could contribute to atmospheric conditions of
high excitation. Yet a handful of otherwise normal stars with T${}_{eff}<$
4600K remain: HD 6833; BD+30°2611; HD 141531; and HD 83212. This extended
survey of helium suggests that only the most luminous Population II stars
($M_{V}$ brighter than $-$1.5), with $T_{eff}\lesssim 4500$ K lack the helium
feature.

###  4.1 Equivalent Widths

The equivalent widths (EW) of the He I $\lambda$10830 absorption are shown in
Figure 8 and 9 as a function of $T_{eff}$ and values are listed in Table 2.
Repeated measurements suggest the error in measuring the equivalent width is
about 5%. The red giants may have an increasing equivalent width with
decreasing effective temperature; this is not unexpected as an extended
expanding atmosphere increases scattering in the line. In the coolest star, HD
104207, the He line is blended with Ti I absorption to shorter wavelengths.
Here, the measurement of the equivalent width is uncertain because of the
blend with both Ti I and Si I. A hint that the absorption is extended and the
equivalent width underestimated in HD 104207 comes from comparison of the
short wavelength side of another Si I line at $\lambda$10843.90 to that of the
line at $\lambda$10827.09 (Figure 6). There may be excess absorption on the
latter line arising from an extended He I profile.

The He I line is surprisingly strong in the RHB stars. In 4 out of 11 RHB
stars, the depth of the line extends 10 to 20 percent below the continuum and
reaches equivalent widths between 0.1 and 0.5Å. These stars also show a
dramatic increase in the helium equivalent width that sets in at
$T_{eff}\gtrsim 5300$K. The values of the equivalent widths are comparable to
those found for Population I stars including binaries which are well-known
X-ray sources (Zarro & Zirin 1986; Sanz-Forcada & Dupree 2008). RHB stars are
not known to be X-ray sources (which would enhance the ionization of He I,
populate the metastable level of He I by recombination and cascade, and create
a stronger line). There are no X-ray sources at the positions of the two stars
with strongest He absorption in the HEASARC Archives
[http://heasarc.gsfc.nasa.gov] indicating that X-ray illumination does not
appear to be present to strengthen the line. The RHB star, BD +17°3248 has a
chromosphere as documented also by the presence of Mg II ultraviolet emission
(Dupree et al. 2007).

The star HD 195636 displays an exceptionally strong He line. Preston (1997)
first noted that this red horizontal branch star is a rapid rotator, which is
confirmed by Carney et al. (2008) as a single star. A strong helium line is
present also in the RHB object, HD 119516 which is not rapidly rotating and
has not been identified as a binary (Carney et al. 2008). Thus rotation does
not appear related to the strength of the helium line, although our sample
consists only of 4 stars. It is interesting to note that these horizontal
branch stars are in a similar helium-burning evolutionary phase as Population
I clump giants. Clump stars, such as the well-studied Hyades giants, exhibit
magnetic activity cycles, ultraviolet emission, and X-rays (Baliunas et al.
1983, 1998).

The increasing strength of the He I line with higher effective
temperature—possibly connected to the development of a hotter
chromosphere—suggests that collisions might be effective in populating via the
forbidden transition from $1s^{2}\ {}^{1}S$ to $1s2s\ ^{3}S$. The strength of
the line in the CH star HD 135148, which has a degenerate companion,
demonstrates that ionization of He by a hot companion, with subsequent
recombination and cascade, can also be important in populating the $1s2s\
^{3}S$ level. Some fraction of stars will have undetected white dwarf
companions; this may be another parameter affecting the strength of the helium
line.

The absorption equivalent widths in this metal deficient sample are generally
lower than found in bright Population I stars of low magnetic activity. In
Figure 10, the absorption equivalent widths for single red giants (Luminosity
classes II-III, III, and III-IV) are taken from the high quality measures in
the sample of O’Brien and Lambert (1986). Many of the spectra show variability
in the equivalent widths, but the He I line is generally stronger in the stars
with roughly solar composition than in the metal-poor sample. A P Cygni
profile may suffer some filling in of the absorption by near-star scattering,
but not all lines display these profiles. The helium abundance, Y, increases
by $\sim$20% between metal-poor and solar models (Girardi et al. 2000), which
is not enough for a factor of $\sim$4 change in the equivalent width. For the
same energy input, the metal-poor chromospheres may be warmer since radiative
losses are less, however that would strengthen the helium line and not weaken
it. Perhaps the helium absorption is enhanced in the Population I stars. They
generally exhibit magnetic activity which leads to X-ray emission that
increases the lower-level population through photoionization followed by
recombination. Semi-empirical models of these chromospheres are needed.

The equivalent width of the $\lambda$10830 line as a function of metallicity
is shown in Figure 11 where no systematic dependence on [Fe/H] appears over
this range of lower [Fe/H]: $-$0.7 to $-$3.0. Red giants of solar metal
abundance show varying strengths of the He I line that tend to cluster between
100 and 200 mÅ, although several stars display values comparable to the metal-
deficient sample.444The stars that are X-ray sources among the Population I
giants show substantially increased strength in the helium lines (Zarro &
Zirin 1986; O’Brien & Lambert 1986; Sanz-Forcada & Dupree 2008), presumably
due to X-ray photoionization, followed by recombination contributing to the
population of the lower ${}^{3}S$ level. It is premature to speculate on the
helium abundance from the line strengths alone without modeling this and other
helium profiles, since they depend on chromospheric conditions. However, if
spectra could be obtained in a globular cluster, providing a larger sample of
similar stars, the relative abundance of helium might be assessed.

###  4.2 Line Profiles

About one-third of all the luminous stars in Figure 7 (AGB and RGB stars
brighter than $M_{V}=0.5$) show helium emission. In most of these stars, the
emission is accompanied by absorption. These classical P Cygni profiles by
their very nature mark an extended outflowing atmosphere. Absorption profiles
without emission can indicate atmospheric dynamics by their asymmetry. The
ratio of the short wavelength extent (at the continuum level) to the long
wavelength extent of an absorption feature relative to the photospheric rest
wavelength gives a measure of the line profile asymmetry. These values are
converted to velocity units and are shown in Figure 12 for the stars without
emission. The value of $B/R$ is given where $B$ denotes the blue (short
wavelength) extent and $R$ the red (long wavelength) extent. The majority of
the helium lines have $B/R>1$ signaling outflowing motions. Values of the
short wavelength extent of the He I absorption are taken as the terminal
velocity ($V_{term}$), and $B/R$ ratios are given in Table 2. It is generally
easy to see from the spectra that helium absorption can extend to the strong
Si I line at 10827.09Å. Such an extension implies an expansion velocity of at
least 90 km s-1. Many stars exhibit higher speeds with absorption evident in
the short wavelength wing of Si I. The outflow velocities measured by the
extent of the short wavelength wing are independent of [Fe/H] as shown in
Figure 13. Metallicity, naturally, is a factor determining the speeds of
radiatively-driven winds. And there is some evidence for dusty winds in OH/IR
sources in the low metallicity Magellanic clouds to have lower speeds as
compared to similar sources in the galactic center (Marshall et al. 2004).
However, the outflow speeds of gaseous winds detected here (presumably driven
by hydrodynamic or magnetic processes) do not depend on the [Fe/H] abundance,
and we conclude that these winds are not radiatively driven.

RHB stars have a convective core and a semi-convective envelope (Castellani et
al. 1971; Schwarzschild 1970) and so conditions exist for the acceleration of
a stellar wind by magnetic processes such as Alfvèn waves. In addition, if
high temperatures are produced in an extended chromosphere, these could
contribute to a thermally driven wind.

The extension of the 10830Å line to shorter wavelengths signals outflow that
in many stars is comparable in value to the escape velocity from the stellar
chromosphere:

$V_{esc}(\rm km\ \rm s^{-1})=620\left(\frac{{\it M/M_{\sun}}}{{\it
R/R_{\sun}}}\right)^{1/2},$ (1)

where $M$ is the stellar mass and $R$ is the distance from the star center to
some region in the chromosphere. We take the 10830Å line to be formed at 2R⋆
in the stellar chromosphere where R⋆ is the stellar photospheric radius. This
estimate does not require a helium model. Detailed calculations concur on the
location of the formation of H$\alpha$ in metal-poor stars. Observations of
higher outflow velocities as well as semi-empirical models confirm that the
10830Å line is formed above the H$\alpha$ core in luminous stars. Our
spherical models for metal-deficient giants (Dupree et al. 1984) have a
chromospheric extent of ’several’ stellar radii in order to produce the
H$\alpha$ line (1.2 R⋆ to 3.6R⋆). The recent (non-LTE, spherical expanding)
models of Mauas et al. (2006) note that the H$\alpha$ core is formed ’about 1
stellar radius above the photosphere’, similar to our spherical models.
Subsequent modeling of H$\alpha$ in M13, M15, M92 giants show the H$\alpha$
cores to be formed at 2R⋆ (Mészáros et al. 2009a). Hence, it appears
reasonable to assume the level of formation of the 10830Å line as 2R⋆. Table 2
contains an estimated stellar radius (Column 10) determined by evaluating the
bolometric correction for each star (Alonso et al. 1999) as a function of
$T_{eff}$ and [Fe/H]. The chromospheric escape velocity for each star is
tabulated in Column 11 of Table 2 using Equation (1). Here we have assumed
masses of a red giant (0.75M⊙), a red horizontal branch star (0.7M⊙), an AGB
star (0.6M⊙), a subgiant branch star (0.8M⊙), and a semi-regular variable
(0.6M⊙). Marked in boldface are values of $V_{term}$ when they are comparable
to or exceed the escape velocity at 2R⋆. These amount to 40% of our sample
where the helium line is detected.

Many of the remaining stars exhibit a short wavelength extension of
$\gtrsim$40 km s-1; this value exceeds the extension of the long wavelength
wing, and signals that outflow of material is present that has not yet reached
escape velocity. Where helium occurs in the luminous stars,
($M_{V}\lesssim-0.2$), a signature of outflow is found in each one. However, a
fraction of the lower luminosity objects do not show this signature,
suggesting that the gas outflow may be variable.555One star, BD+17°3248
($M_{V}=0.65$), was observed 3 years previously (Smith et al. 2004) and the
helium profile has not changed. Speeds greater than $\sim$ 12 km s-1 are
generally supersonic in fully ionized metal-deficient ([Fe/H]=$-$2) plasma at
chromospheric temperatures of 104K.

Three of the target stars (BD +17°3248, HD 122956, and HD 126587) have high
resolution Hubble Space Telescope spectra available of the Mg II line at
$\lambda$2800 (Dupree et al. 2007). Asymmetries of the line emissions indicate
motions in the chromosphere, and these were found in stars brighter than
MV=$-$0.8. Only the most luminous of the 3, HD 122956, shows a Mg II emission
asymmetry indicating outflow (short-wavelength emission peak less than the
long-wavelength emission peak). HD 122956 has a high value of the helium
terminal velocity, 110 km s-1 which exceeds the chromospheric escape velocity
of 78 km s-1. The RHB star, BD +17°3248, shows outflow in helium at an
intermediate velocity, whereas the helium line in HD 126587 appears symmetric.
However, models (Dupree et al. 1992a) suggest that Mg II is formed at lower
levels than the 10830Å line in a metal-deficient chromosphere.666Other cool
stars have observational signatures of this separation: Cepheids (Sasselov &
Lester 1994), a T Tauri star (Dupree et al. 2005), and the Sun (Avrett 1992).
Thus it is not surprising to find differing dynamical signatures in these line
diagnostics, in addition to possible time variations. There may be a
similarity here between the well-known changing asymmetries of H$\alpha$
emission wings in metal deficient giants (Smith & Dupree 1988; Cacciari et al.
2004; Mészáros et al. 2008, 2009b) and the Mg II emission.

The star HD 135148 deserves special mention. This RGB object is classified as
a CH star and Carney et al. (2003) obtained an orbital period of 1411 days for
the spectroscopic binary. Emission in the P Cygni profile arises from an
extended scattering atmosphere and the absorption extends to $\sim$$-$115 km
s-1. This value exceeds the escape velocity from the chromosphere,
$V_{esc}$=67 km s-1, where helium originates. Thus after the initial transfer
of material to the secondary star in the system, a substantial wind remains
from the cool star that is presently visible.

###  4.3 Estimate of the Mass Loss Rate

A rough estimate of the mass loss rate implied by the helium absorption in the
wind can be derived from the Sobolev optical depth. In an expanding
atmosphere, a photon emitted from the photosphere (or the line itself) can be
absorbed and then scattered when the absorption coefficient is “aligned” with
the photon. In this situation, a sufficient number of atoms occur at the
correct velocity to absorb and scatter photons. In a stellar wind, a narrow
interaction region will be present that depends on the velocity gradient in
the wind, and the width and strength of the line absorption coefficient. The
Sobolev approximation defines the interaction region to be very narrow for
simplification of the transfer equation, and this causes the absorption
parameters to be related only to local conditions (Lamers & Cassinelli 1999).
This approximation assumes that the density and velocity gradient do not
change significantly over the absorbing/scattering region.

The line optical depth at frequency, $\nu$, at star center, is given by

$\tau_{\nu}=\int_{0}^{\infty}\kappa_{\nu}(z)\rho(z)dz$ (2)

along the radial direction, $z$, where $\kappa_{\nu}$(cm2 g-1) is the line
absorption coefficient, and $\rho$(g cm-3) is the mass density in the lower
level of the transition. Taking the line profile function as a delta-function
(the Sobolev approximation) and inserting values for $\kappa_{\nu}$ [cf.
Equation (8.51) of Lamers & Cassinelli (1999) or Equation (8.8) of Hartmann
(1998)], we write the Sobolev optical depth, $\tau_{S}$ as:

$\tau_{S}=\frac{\pi e^{2}}{mc}\times
f\times\lambda_{0}\times\frac{N_{1}}{(dV/dz)}.$ (3)

where $dV/dz$ is the velocity gradient in the scattering region and $N_{1}$ is
the population in the lower level of the absorption line (the ${}^{3}S$ level
of He I). Conservation of mass gives

$\dot{M}=4\pi R^{2}V\mu m_{H}N_{H}(V)$ (4)

where $R$ is the radial distance (in units of $R_{\sun}$) at which the wind
has a velocity $V$(km s-1). NH is the hydrogen density at $V$, $\mu$ is the
mass per hydrogen nucleus, and mH is the mass of the hydrogen atom. Then
substituting for $N_{H}$ into Equation (4) [rewriting
$N_{H}=N_{1}\times(N_{H}/N_{1})$], and replacing the value of $N_{1}$ from the
expression for $\tau_{S}$ above, we find

$\dot{M}=\frac{4\pi R^{2}V\mu
m_{H}}{N_{1}/N_{H}}\times\frac{\tau_{S}}{\frac{\pi
e^{2}}{mc}f\lambda_{0}}\times\frac{dV}{dz}.$ (5)

We assume that $dV/dz\sim\Delta V/\Delta R=V/(R-R_{\star})$ where $\Delta V$
is the change in wind velocity, ($R-R_{\star}$) is the distance over which the
speed changes from zero at the stellar photosphere to a value of $V$ at
distance $R$. In our estimate, we adopt $R=2R_{\star}$ since the 10830Å line
is formed in the chromosphere (see discussion in Section 4.2) and $R$ is
measured from the center of the star. Then with $\mu=1.4$, and
$m_{H}=1.67\times 10^{-24}$ g, we have:

$\dot{M}\ {(\rm M_{\sun}\ \rm yr^{-1})}=\frac{1.22\times
10^{-18}\tau_{S}(R/R_{\sun})^{2}V^{2}}{(N_{1}/N_{H})\times
f\times\lambda(\AA)\times(R_{\star}/R_{\sun})}$ (6)

where $N_{1}/N_{H}$ is the ratio of the population in the lower ${}^{3}S$
level of the helium transition to the total hydrogen density. To evaluate the
mass loss rate from Equation (6) at a distance of 1 $R_{\star}$ above the
photosphere, we set $\tau_{S}$ =1 and the oscillator strength, $f=0.54$ for
the $\lambda$10830 multiplet. With these values, the mass loss rate becomes,

$\dot{M}\ (\rm M_{\sun}\ \rm yr^{-1})=\frac{8.37\times 10^{-22}{\it
R_{\star}V^{2}_{term}}}{{\it(N_{1}/N_{H})}}$ (7)

where $R_{\star}$ is the stellar (photospheric) radius (in units of
$R_{\sun}$), and $V_{term}$ is the observed terminal velocity (km s-1) in the
helium line.

Now, the value of $N_{1}/N_{H}$ can be estimated using our semi-empirical non-
LTE models (Avrett & Loeser 2008) of cool star chromospheres (see Appendix A).
The semi-empirical models suggest that an upper limit of the population ratio,
$N_{1}/N_{H}$ for typical line strengths found in the targets reported here
(line depth $\sim$0.9) corresponds to 6.3$\times$10-8 where
$N_{He}/N_{H}=0.1$. As a lower limit, we take the value derived from the solar
model: 1.0$\times$10-8. The mass loss rates are estimated using values for
$V_{term}$ (where $V_{term}>45$km s-1) and $R_{\star}$ contained in Table 2,
and are shown in Fig. 14. Two rates are plotted for each star corresponding to
the upper and lower limit on $N_{1}/N_{H}$. The mass loss rate generally
increases with increasing stellar bolometric magnitude. The uncertainty in
these estimates arises predominantly from the population of the lower
${}^{3}S$ level of the $\lambda$10830 transition. A discussion of the error
estimate is given in Appendix B. Equation (7) does not apply to the
exceptionally deep P Cygni profile of HD 135148 such as that shown in Fig. 5
because it likely overestimates the mass loss rate.

Undoubtedly there is variation in the gas mass loss rate for stars on the RGB
- and probably all stars considered here. The solar mass loss rate changes
over the solar cycle by about a factor of 1.5 (Wang 1998). Mészáros et al.
(2009a) found variations ranging between a factor of 2 to 6 in the mass loss
rate of individual globular cluster red giants. Accompanying the changes in
the mass loss rate may be variations in the size of the outflow velocities.
However several observations suggest that outflows occur continually as stars
evolve through the upper part of the RGB. The helium line profiles reported
here overwhelmingly display a signature of outflowing gas in stars brighter
than $M_{V}\sim 0$. H$\alpha$ line cores show generally increasing outflowing
velocities with luminosity (Mészáros et al. 2008, 2009b). Even though some
measured outflow velocities are less than the escape velocity, conservation of
mass suggests that the velocities will yield meaningful mass loss rates. We
emphasize the difference between the presence of mass loss from red giants
determined from diagnostics of the gas (from He I 10830Å and the H$\alpha$
line) and that derived from infrared detections of circumstellar dust.
Evidence suggests that dust formation is an episodic process (Origlia et al.
2007; Mészáros 2008), whereas the velocity measurements of the gas indicate
continuous outflow of material.

With regard to the presence of a wind directly indicated by the helium line
profile, a fraction of the stars in Table 2 and 3 show velocities exceeding
the chromospheric escape velocity. There are several possible interpretations
of these measurements: (1) Only some fraction of the stars have outflows that
develop into winds; (2) All stars develop winds for some fraction of their
evolutionary phase brighter than $M_{V}\sim 0$. (3) The outflow velocities
vary, probably along with the mass loss rates, and we need diagnostics formed
higher in the atmospheres to attempt the detection of escape velocities. The
absence of inflowing velocities both in H-$\alpha$ and the helium line would
suggest that the last option appears to be the most likely.

###  4.4 Evolutionary Effects of Mass Loss

Stars on the red giant branch exhibit mass loss rates that increase with
luminosity. At a magnitude comparable to the horizontal branch, $\dot{M}\sim
3.2\times 10^{-9}$ to $3.2\times 10^{-8}$ $M_{\sun}\ yr^{-1}$. At higher
luminosities on the RGB, the mass loss rates increase by about a factor of 2,
viz., $6.3\times 10^{-9}$ to $6.3\times 10^{-8}$$M_{\sun}\ yr^{-1}$. A low-
mass star spends about 50 Myr in the RGB phase, thus, taking the geometric
mean of the minimum and maximum values, ($4.5\times 10^{-9}$$M_{\sun}\
yr^{-1}$), the total mass lost as indicated by the helium line would amount to
$\sim$ 0.2 $M_{\sun}$. This is the amount generally demanded by stellar
evolution considerations which range from $\sim$0.15 to 0.22$M_{\odot}$ (Rood
1973; Lee et al. 1994; Caloi & D’Antona 2008; Dotter 2008). Moreover, this
mass loss rate is consistent with values derived from the H$\alpha$ profiles
of globular cluster giants (Mauas et al. 2006; Mészáros et al. 2009a).

Stars on the AGB have estimated mass loss rates of 0.25$-$1.6$\times$ 10-8
$M_{\sun}\ yr^{-1}$. For a 20 Myr lifetime on the AGB, additional mass loss of
$\sim$ 0.1M☉ could result for the AGB objects here (assuming a mean mass loss
of 6.3 $\times$ 10-9$M_{\sun}\ yr^{-1}$). Groenewegen & deJong (1993) find
about 0.16$-$0.2 $M_{\odot}$ is lost on the AGB. The sample of AGB stars
included here has low luminosities ($M_{bol}\approx$ 0 to $-$1.9), and it is
likely that the mass loss increases with luminosity which would increase the
empirical total mass loss.

Minimum and maximum values of the mass loss rate for the RHB stars in our
sample range from $4.5\times 10^{-10}$ to $2.2\times 10^{-8}$ $M_{\sun}\
yr^{-1}$(Fig. 14). Taking a median of $3\times 10^{-9}$$M_{\sun}\ yr^{-1}$,
and a lifetime of 75 Myr, this rate implies a total mass loss of
0.2$M_{\odot}$. Note that only half of the sample of RHB stars has an
asymmetric profile indicating outflow velocities, suggesting that the outflow
and hence the mass loss rate may be lower. Some evolutionary models have
considered mass loss on the RHB. Yong et al. (2000) and earlier, Demarque &
Eder (1985) evaluated models of horizontal branch stars with the ad hoc
assumption of mass loss to test its effects, and concluded that mass loss
rates between 10-10 and 10-9 $M_{\sun}\ yr^{-1}$ could produce the observed
extended blue HB. Koopmann et al. (1994) set an upper limit to the HB mass
loss rate of 10-9 $M_{\sun}\ yr^{-1}$, based on calculations for M4. These
values are not substantially discrepant from the mass loss rate inferred from
the He I line profiles.

Vink and Cassisi (2002) evaluated the effect of radiatively driven winds in
horizontal branch stars, but found the mass loss rates too low for
evolutionary effects. However, winds can be driven in several other ways which
is likely here, given the presence of convection layers in these stars
(Castellani et al. 1971; Schwarzschild 1970). It is worth noting that even
models of RGB and AGB stars currently consider sub-surface dynamo activity
(Busso et al. 2007, Nordhaus et al. 2008) which could lead to winds driven by
magnetic processes.

Carney et al. (2008) noted that red horizontal branch stars in the field are
rotating slower than expected considering the rotation observed in their
predecessors on the red giant branch. Angular momentum carried away by the
stellar winds could offer an explanation of this discrepancy. Here we estimate
the torque required to spin down a red giant with the mass loss rates
discussed above. To be effective, magnetic fields must be present and we
assume that the moment of inertia of the star does not change. The torque
required, $\tau_{\star}$ (dyne cm) under these conditions is
$\tau_{\star}=(I_{o}\omega_{o})/\Delta t$ where $I_{o}$ is the stellar moment
of inertia, $\omega_{o}$, the initial angular velocity, and $\Delta t$ the
time required to completely stop the rotation of the red giant. Taking, a
tangential rotation velocity of 2 km s-1 [from the Carney et al. (2008)
measures, and setting $sin\ i=0.3$] and R⋆=50R☉ for a 0.8M☉ red giant star, we
find that a torque of 7$\times$1035 dyne cm is required to spin down the star
in 20 Myr. Matt & Pudritz (2008) evaluated the torque created on a star by
winds of various mass loss rates in the presence of a stellar magnetic field
with several configurations and strengths. A simple parameterization of the
quantity: $(B_{\star}R_{\star})^{2}/(\dot{M}V_{esc})$ can predict the ’lever
arm’ to evaluate the torque exerted by the wind on the star. For a red giant
with the parameters above, a dipole magnetic field of 100 gauss, and
$\dot{M}$= 4.5 $\times$10-9 $M_{\sun}\ yr^{-1}$, the predicted wind torque
amounts to 3 $\times$ 1036 dyne cm, a value that exceeds the amount required
for spindown, making winds a plausible explanation of the low velocities of
RHB stars. However we caution that a change in the moment of inertia could
affect these results as well as the magnetic field strength and
configuration.777Y. C. Kim made available his calculations of the moment of
inertia for a 1 M☉ star with [Fe/H]=$-$0.25 as it evolves up the red giant
branch. These show that the moment of inertia increases by a factor of 6.5,
based on internal structural changes including an increase in radius, as the
star evolves from an effective temperature of 4500K to 4000K. Such a factor
would appear to allow spin-down to occur.

These results can be used to calculate a total budget of mass loss for post-
main sequence evolution. Taking values of $\dot{M}$ estimated above, we find a
star would lose 0.2 M☉ on the RGB, 0.1–0.2 M☉ on the RHB, and 0.1–0.2 M☉ on
the AGB and as a planetary nebulae (Bianchi et al. 1995), totaling 0.4–0.6 M☉.
This total is in harmony with recent studies of globular clusters that suggest
anywhere from 0.5 to over 1 $M_{\odot}$ is lost between the main sequence
turnoff and the white dwarf cooling sequence (Moehler et al. 2004; Hansen et
al. 2007; Richer et al. 2008). Our budget should not be strictly taken to
apply to current turnoff stars of mass $\sim 0.8$ M☉, since many of the white
dwarfs observed in clusters are remnants of stars of higher initial masses,
where more mass loss is demanded in post-main sequence evolution. In addition,
the Population II field giants that make up our sample potentially had a range
of initial masses, and thus some may have mass loss rates higher than expected
for globular clusters.

###  4.5 The Fate of Wind Material

If red giants in globular clusters have winds similar to those detected among
Population II field giants, it is worth considering whether the wind material
would have enough energy to escape not only from the stars themselves but also
from the parent cluster. Escape velocities from the cores of Milky Way
globular clusters have been evaluated by McLaughlin and van der Marel (2005).
Central escape velocities vary from $\sim$ 2 km s-1 (AM1, Pal 5 and Pal 14) to
$\sim$ 90 km s-1 for massive clusters (NGC 6388 and NGC 6441). A majority of
the stars brighter than MV = 0 in our sample exhibit terminal velocities in
excess of the stellar escape velocity from the chromosphere. These speeds,
ranging from 90 to 170 km s-1 also exceed many of the escape speeds from
clusters too. Since the cluster escape speed corresponds to the speed
necessary to remove material from the core, the value will obviously decrease
with distance from the core. One could envision a scenario in which material
does not have sufficient energy to escape the cluster core, but it could
escape if arising from a star located further out in the cluster.

Several authors have considered in detail the dynamics of wind material
deposited in the intracluster medium by the stellar population. VandenBerg and
Faulkner (1977) used hydrodynamic equations to construct time dependent models
of gas flow from a cluster assuming the stellar wind has a velocity $\sim$20
km s-1. Depending upon the initial assumptions of the energy available, they
found that both outflows and inflows of material could occur, the latter as a
result of radiative cooling. In some models material was retained in the
cluster core. Smith (1999) concluded that additional energy could be injected
into the gas by solar-like winds originating from cluster dwarf stars and this
would be sufficient to establish an outflow of material from a cluster even if
red giant winds are slow. If the giant winds are fast, comparable to the solar
wind with V $\sim$ 450 km s-1, steady state outflows would result, even from
clusters with the highest values of escape velocity (Faulkner & Freeman 1977).

The results presented here from helium lines suggest that the winds from red
giants and red horizontal branch stars can be much faster than generally
assumed, and these would naturally lead to a steady-state outflow.

Recent observations of the fast solar wind suggest that interactions between
colliding winds within the cluster itself may also be of some consequence.
Insight can be drawn from the Sun where spacecraft have located the
termination shock and characterized its physical parameters (Richardson et al.
2008). The shock occurs $\sim$100 AU distant from the Sun, and the wind speed
at that point is comparable to the speed in the corona, namely 400 km s-1.
While the solar wind is hot and driven by a combination of gas and wave
pressure, a cool wind can be driven by wave pressure alone (Cranmer 2008).
However, convective envelopes in red horizontal branch stars, and recent
conjectures of dynamo activity in RGB and AGB stars would allow MHD processes
to occur as well (Busso et al. 2007; Nordhaus et al. 2008). Thus, it appears
possible that the stellar winds will not be substantially decelerated up to
the termination shock, but maintain the values indicated by the helium line.

At large distances from the star, the pressure of the stellar wind will
eventually be balanced by the opposing pressure (both gas and magnetic)
presented by the surrounding intercluster medium, $P_{icm}$. When these
pressures are equal, a termination shock occurs at some distance, $R_{TS}$,
viz.:

$R^{2}_{TS}=\frac{\dot{M}_{\star}\times V_{wind}}{4\pi P_{icm}}.$ (8)

Rewriting this using astronomical units,

$R_{TS}(AU)=48.4\times\sqrt{\frac{\dot{M}_{\star}(M_{\sun}\ yr^{-1})\times
V_{wind}(km\ s^{-1})}{P_{icm}(dyne\ cm^{-2})}}$ (9)

and assuming $\dot{M}=$2$\times$10-9 M☉ yr-1 for an average RGB star, Vwind=
100 km s-1, and a (total) interstellar pressure of 4.0$\times$10-13 dyne cm-2
for a distance of 1.5 kpc above the Galactic plane (Cox 2005), we find the
transition shock distance is 34,000 AU or 0.2 pc from the star. The central
luminosity densities for globular clusters have median values of $\sim$ 3000
$L_{\sun}pc^{-3}$ (Harris 1996), and with a typical star of 0.1 to 0.2
$L_{\sun}$, the median stellar density in the core is greater than 104 stars
pc-3, giving an average separation between stars of less than 9500 AU -
smaller than the termination shock distance. While uncertainty exists in both
the pressure outside and within clusters888The pressure in the galactic halo
is not firmly known. O VI absorption suggests warm (3 $\times$105K) extended
low density regions are present at high galactic latitudes (Dixon and Sankrit
2008) with thermal pressures of 0.7 to 1 $\times$ 10-12 dyne cm-2. Electron
densities in the intercluster material of 47 Tuc derived from pulsar
dispersion measures (Freire et al. 2001) indicate $n_{e}=0.067\pm 0.015\
cm^{-3}$. They concluded ionized material was dominant, where for a
temperature of 104K, the thermal pressure equals 9$\times$10-14 dyne cm-2. van
Loon et al (2006) detected 0.3$M_{\sun}$ of material in the core of M15 using
the Arecibo telescope. Assuming that this material is evenly distributed in
the beam we find a hydrogen density of 2.29$\times$10-2 cm-3 for the diffuse
gas in the core. With a temperature of 100K, the pressure will be
3.2$\times$10-16 dynes cm-2. Faulkner and Freeman (1977) constructed time-
independent gas flow models for tightly bound clusters. These models suggest
pressures at the tidal radius ranging from 0.1 to 5.5 $\times$10-16 dyne cm-2.
At the sonic point in the flow which lies in the cluster interior where the
stellar density is down by a factor of $\sim$100 from the core density, the
gas pressures are much higher, ranging from 5$\times$10-15 to 3.5$\times$10-12
dynes cm-2. Our estimate of the transition shock distance varies inversely as
the square root of the external gas pressure, so that an order of magnitude
change in the gas pressure, changes the distance by a factor $\sim$ 3., these
estimates suggest that the central regions of clusters may well be filled with
expanding warm material, enveloping the surrounding stars. Since the majority
of cluster stars are still on the main sequence with winds of $\dot{M}\sim
2\times 10^{-14}$ $M_{\sun}\ yr^{-1}$ and $V_{wind}\sim 400$ km s-1 (adopting
solar parameters), these winds can not balance the pressure of the more
massive winds arising from the luminous stars, raising the possibility that
the dwarf winds are in fact smothered by the giant winds, allowing for
pollution of the surface layers of the dwarfs. As the stellar density
decreases towards the cluster edges, the effect of smothered winds and
resultant surface pollution would decrease, possibly leading to spatially
dependent self-pollution within a cluster. Even beyond the termination shock
crossing, the wind speed decreases by about a factor of 2 (based on the solar
termination shock measured by Richardson et al. 2008), but the temperature and
density increase. In many cases, even the decreased speed will allow escape of
these fast winds from the cluster.

Central escape velocities for the most massive clusters can reach $\sim$90 km
s-1 (McLaughlin & van der Marel 2005). While the helium line profile suggests
that some objects possess these high velocities, others do not (see Figure
13). The extent of velocity variability in the helium line is unknown at
present. The velocity of a red giant wind at levels higher than the formation
region of the helium line may (or may not) reach escape velocities.
Calculations of line forming regions for strong optical, IR, and near UV lines
have been made for metal deficient red giants (Dupree et al. 1992a; Mauas et
al. 2006; Mészáros et al. 2009a) but only through the low chromosphere and not
for higher levels of the atmosphere. In the high chromosphere, the most
straightforward diagnostics of winds lie in the ultraviolet region of the
spectrum, where Lyman-$\alpha$ and resonance lines of C II ($\lambda$1335) and
Si II ($\lambda$1800) might be observable.999UV and Far-uv spectra obtained
with IUE, HST, and FUSE of the brightest of the metal-deficient targets in the
field, HD 6833, are weak and can not definitively reveal emission. One
advantage of metal-poor targets is that their high velocities move the spectra
away from local interstellar absorption which could compromise the resonance
line profiles. Infrared transitions of hydrogen such as members of the
Paschen, Brackett, and Pfund series arise from higher levels of hydrogen and
hence are formed deeper in the atmosphere and can not help in detecting a
wind.

##  5 Conclusions

The He I 10830Å transition maps atmospheric dynamics to higher levels than
optical or near-uv diagnostics. The profile of this near-infrared line in
metal-poor stars gives evidence for outflow in most of the RGB, AGB, and RHB
stars showing helium. In many objects, the speeds are comparable to the escape
velocities from both the stellar chromosphere and globular clusters. The fact
that all the luminous stars with helium absorption exhibit expanding
chromospheres suggests that the mass outflows as traced by gas occurs
continually, with most probably variable mass loss rates.

Our estimate suggests that the mass loss observed directly on the RGB will
provide the requisite amount needed by stellar evolution calculations. Mass
loss detected in RHB stars appears at a rate sufficient to cause extension of
the horizontal branch. It will be very useful to obtain ultraviolet spectra of
some of these stars to identify higher temperature plasma, and to track the
acceleration in the chromosphere. Better estimates of mass loss rates can be
obtained with semi-empirical modeling of the line profiles.

These results demonstrate that chromospheric material has sufficient speed to
escape these stars and become a stellar wind. If the star were in a globular
cluster, the high-speed wind continues with little diminution, filling the
cluster with expanding warm gas. In the process, red giant winds could smother
the substantially less massive winds from dwarf stars possibly allowing for
surface pollution. Beyond the termination shock, velocities decrease but could
still escape the cluster. Thus fast winds observed in the helium line offer a
straightforward way to understand the absence of intracluster material in
globular clusters.

We are grateful to Steve Cranmer and Aad van Ballegooijen for insight into the
solar wind. Gene Avrett kindly made the helium calculations for the Sun
available before publication. We thank Y. C. Kim for providing us with his
pre-publication results for the evolution of the moment of inertia of red
giants. This research has made use of NASA’s Astrophysics Data System Abstract
Service and SIMBAD database (CDS, Strasbourg, France). 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: Keck 2 (NIRSPEC)

##  Appendix A Estimation of Level Population

The level population for the ${}^{3}S$ level of He I, n(${}^{3}S$), is
estimated from non-LTE calculations using semi-empirical models for different
stars. A detailed discussion of helium excitation processes in the solar
atmosphere is given in Andretta & Jones (1997) and is not reviewed here. Our
estimates of the population levels derive from the PANDORA code (Avrett &
Loeser 2003) which was used in either plane-parallel or spherical form, with
an expanding wind in most models. A 13-level helium atom plus continuum was
generally used; some models had a 5-level helium atom plus continuum to
expedite calculation. The velocity field was introduced explicitly into the
source functions. Solutions for hydrogen populations and ionization are
iterated first, and then they are followed by similar iterations for helium.
The highest value n(${}^{3}S$)/n($He_{tot}$) in each model is shown in Fig. 15
as a function of the maximum depth of absorption in the helium 10830Å line.
The solar model represents the quiet sun for a static plane-parallel
atmosphere which is described in detail elsewhere (Avrett & Loeser 2008), and
E. Avrett kindly made an advance copy of the helium results available for use
here. High It may well be that the mass loss rates are variable. For instance
only one-half of the RHB stars exhibit outflow. And recent chromospheric
models calculated for varying H$\alpha$ profiles in globular cluster red
giants, reveal changes in mass loss rate by a factor of 6 (Meszaros et al.
2009a). Thus the total mass lost would seem to intrinsically span a range of
values. luminosity stars (similar to giants and supergiants) were also
calculated in plane-parallel and spherical models both in a static atmosphere
and also with an assumed mass outflow. The models were of cool stars (such as
$\beta$ Dra, $\alpha$ Aqr, and $\alpha$ Boo) with effective temperatures
approximately solar. The photospheric model has little or no influence on the
helium line profile since it is formed totally in the chromosphere. Results
for Mg II and helium lines in a supergiant were described elsewhere (Dupree et
al. 1992b). A cool dwarf star with an extended chromosphere (TW Hya) producing
a P Cygni profile was also modeled in helium, with a spherical approximation
and a substantial high velocity wind (Dupree et al. 2008).

Figure 15 shows the maximum depth of the He I 10830Å as a function of the
population ratio n(${}^{3}S$)/n($He_{tot}$) of helium. The three plane
parallel models (for the Sun and the supergiant and giant stars marked with
vertical lines in Figure 15) show generally lower values of the population
ratio. As one might expect, there is a correlation between the depth of the
line and the level population. Higher level populations result in increasing
the line depth. The models with high outflow velocities in the chromosphere
(110$-$200 km s-1) at the atmospheric level where the ${}^{3}S$ population
maximizes, create deeper absorption profiles, but the level populations are
remarkably similar. One of our target stars, HD 135148 has an extreme 10830Å
line depth ($\sim$0.1) likely caused by radiation from the hot companion
contributing to the photoionization-recombination processes populating the
lower ${}^{3}S$ level. Other stars observed in this paper have absorption
depths relative to the continuum between 0.8 and $\sim$0.95. Fig. 15 shows
that such absorption depths occur for a population ratio of log
n(${}^{3}S$)/n($He_{tot}$) $<\ -$6.2. This provides our estimate of an upper
limit to the population ratio, which in turn translates into a lower limit to
the inferred mass loss rate for the discussion in Section 4.3.

These models and the assumptions of the Sobolev approximation contain
uncertainties. The atomic physics (collisional excitation rates,
photoexcitation and photoionization cross-sections, and radiative and
dielectronic recombination rates) are continually updated, and these
calculations of populations were made more than a year ago. Only the giant
model for $\alpha$ Boo was constructed using a metal-poor abundance set. For
the same chromospheric input energy, one might expect the chromospheric
temperatures to be higher in a metal-poor environment because the radiative
losses are less. However the Alpha Boo (giant) model in which the abundances
are a factor of 3 less than solar appears to give results in harmony with
those from solar abundance models. Moreover, these are ’semi-empirical’
models, generally constructed to fit observed line profiles, so that the
required temperature/density structure accommodates a non-solar metallicity.

Our use of the Sobolev approximation introduces assumptions as well. We have
adopted a conservative value for dV/dz. By setting this gradient equal to
$V/(R-R_{\star})$ in section 4.3, our estimate basically averages the entire
velocity increase from 0 km s-1 to the observed He I outflow velocity over the
entire 1 stellar radius of the chromosphere between the surface and the radius
of He I formation. The wind might become more accelerated with increasing
distance from the photosphere, [and acceleration has been observed in red
giants in globular clusters (Mészáros et al. 2009b)], or else be
preferentially accelerated at higher altitudes. This would cause our estimate
of dV/dz (and hence the mass loss rate) to be an underestimate. We have taken
the region of helium line formation as 1 stellar radius above the photosphere,
based on several models of the level of formation of the H$\alpha$ line
placing it at 2R⋆ in globular cluster red giants (Mauas et al. 2006; Mészaáros
et al. 2009a). However, helium is formed above the H$\alpha$ line and so this
could also lead to an underestimate of the mass loss rate. We have set
$\tau_{S}$ equal to 1.0 in the Sobolev approximation, following other authors
(Hartmann 1998). It probably would be less in the chromosphere in a fully non-
LTE calculation, and thus decrease the mass loss rate, but other atmospheric
parameters would undoubtedly change.

As a next step, calculations of the level populations and the line profiles
should be carried out in non-LTE, with spherical coordinates, and include
velocity fields for a grid of temperatures and values of [Fe/H]. It would be
optimum also to have other chromospheric lines than helium, such as H$\alpha$,
Ca II K, and Mg II to constrain the chromospheric structure. While the
H$\alpha$ profiles are known to vary, currently we are ignorant of possible
changes in Ca II, Mg II, and the He I 10830Å lines. Of course, such diagnostic
lines must be acquired with a variety of instruments on the ground and from
space, and realistically may be difficult to achieve - much less achieve
simultaneously. No measure exists of X-rays from these sources; such a
measurement could also provide constraints on any X-ray illumination of the
chromosphere.

It is reassuring that the mass loss rates for stars on the red giant branch
span values of $\approx 3\times 10^{-10}$ $M_{\sun}\ yr^{-1}$ to $6.3\times
10^{-8}$$M_{\sun}\ yr^{-1}$, and these values are congruent with those
determined by non-LTE spherical models of the H$\alpha$ line. Mauas et al.
(2006) find values ranging from $1.1\times 10^{-10}$ to $3.9\times 10^{-9}$
$M_{\sun}\ yr^{-1}$ for 5 red giants in NGC 2808. Mészáros et al. (2009a) find
similar values of the mass loss rate: $5.7\times 10^{-10}$ to $4.8\times
10^{-9}$ $M_{\sun}\ yr^{-1}$ for 15 red giants in 3 globular clusters (M13,
M15, and M92) from the H$\alpha$ line too. At the highest luminosity, the
values from helium derived here appear to be higher than the H$\alpha$
modeling, but in agreement with the classical ’Reimers’ value (Reimers 1975).
There may well be wind variability too, and this could cause changes in the
mass loss rate by perhaps a factor of 6, based on the H$\alpha$ variability
(Mészáros et al. 2009b). Further work would be beneficial to establish mass
loss rates.

##  Appendix B Propagation of Errors in the Mass Loss Rate

We estimate the propagated rms error $\sigma$($\dot{M}$) in the mass loss rate
by evaluating the contribution of 5 variables to $\dot{M}$: R, V, $N_{rel}$=
$N_{1}/N_{H}$, $\tau_{S}$, and $D_{V}$= $dV/dz$. Here, as defined in Section
4.3, $R$ is the distance from the center of the star with radius $R_{\star}$,
$V$ is the expansion velocity, $N_{rel}$ is the population of the lower level
of the He I 10830Å line relative to hydrogen, $\tau_{S}$ is the Sobolev
optical depth, and $D_{V}$ is the velocity gradient. The error is given by:

$\sigma^{2}(\dot{M})=\sum_{X=R,V,N_{rel},\tau_{S},D_{V}}\left(\frac{\partial\dot{M}}{\partial
X}\right)^{2}\sigma^{2}(X).$ (B1)

By evaluating $\left(\frac{\sigma(\dot{M})}{\dot{M}}\right)^{2}$,
$\sigma(\dot{M})$ can be estimated as a function of $\dot{M}$ to assess the
contribution of each quantity to the error. Writing Equation (B1) in terms of
the variables gives,

$\displaystyle\left(\frac{\sigma(\dot{M})}{\dot{M}}\right)^{2}$
$\displaystyle=$
$\displaystyle\frac{1}{\dot{M}^{2}}\left(\frac{\partial\dot{M}}{\partial
N_{rel}}\right)^{2}\sigma^{2}(N_{rel})+\frac{1}{\dot{M}^{2}}\left(\frac{\partial\dot{M}}{\partial\tau_{S}}\right)^{2}\sigma^{2}(\tau_{S})+\frac{1}{\dot{M}^{2}}\left(\frac{\partial\dot{M}}{\partial
V}\right)^{2}\sigma^{2}(V)$ (B2) $\displaystyle+$
$\displaystyle\frac{1}{\dot{M}^{2}}\left(\frac{\partial\dot{M}}{\partial
R}\right)^{2}\sigma^{2}(R)+\frac{1}{\dot{M}^{2}}\left(\frac{\partial\dot{M}}{\partial
D_{V}}\right)^{2}\sigma^{2}(D_{V})$

From Equation (5) in Section 4.3,

$\dot{M}=K\frac{R^{2}V\tau_{S}D_{V}}{N_{rel}}$ (B3)

where $K$ is a constant composed of atomic and scaling parameters. Taking the
partial derivatives of these variables, and reinserting the expression for
$\dot{M}$, we find,

Substituting these quantities into Equation (B2) yields,

$\left(\frac{\sigma(\dot{M})}{\dot{M}}\right)^{2}=\left(\frac{\sigma(N_{rel})}{N_{rel}}\right)^{2}+\left(\frac{\sigma(\tau_{S})}{\tau_{S}}\right)^{2}+\left(\frac{\sigma(V)}{V}\right)^{2}+\left(\frac{2\sigma(R)}{R}\right)^{2}+\left(\frac{\sigma(D_{V})}{D_{V}}\right)^{2}$
(B4)

Now supposing that $\frac{\sigma(N_{rel})}{N_{rel}}=3$,
$\frac{\sigma(\tau_{S})}{\tau_{S}}=1$, $\frac{\sigma(V)}{V}=\frac{1}{2}$,
$\frac{\sigma(R)}{R}=\frac{1}{2}$, and $\frac{\sigma(D_{V})}{D_{V}}=1$ and
substituting these values into Equation (B4), we find

$\left(\frac{\sigma(\dot{M})}{\dot{M}}\right)^{2}=9+1+\frac{1}{4}+1+1$ (B5)

so that the largest contribution to the error arises from the uncertainty in
$N_{1}$, and

$\sigma({\dot{M}})=3.5\dot{M}.$ (B6)

Thus, in this Sobolev approximation, the uncertainty in the mass loss rate
principally depends on the population of the lower level of the He I 10830Å
transition.

## References

  * (1) Alonso, A., Arribas, S., & Martínez-Roger, C. 1999, A&AS, 140, 261
  * (2) Anderson, J. 2002, ASP Conf Ser. 265, ed. F. van Leeuwen, J.D. Hughes, & G. Piotto (San Francisco: ASP), 87
  * (3) Andretta, V., & Jones, H. P. 1997, ApJ, 489, 375
  * (4) Andrievsky, S. M., Korotin, S. A., & Martin, P. 2007, A&A, 464, 709
  * Anthony-Twarog & Twarog (1994) Anthony-Twarog, B. J., & Twarog, B. A. 1994, AJ, 107, 1577
  * Anthony-Twarog & Twarog (1998) Anthony-Twarog, B. J., & Twarog, B. A. 1998, AJ, 116, 1922
  * (7) Avrett, E. H. 1992, in Proceedings of the Workshop on the Solar Electromagnetic Radiation Study for Solar Cycle 22, ed. R. F. Donnelly (Virginia: National Technical Information Service), 20
  * Avrett & Loeser (2003) Avrett, E. H., & Loeser, R. 2003, in IAU Symp 210, Modeling of Stellar Atmospheres, ed. N. Piskunov, W. W. Weiss, & D. F. Gray (San Francisco: ASP), A-21
  * (9) Avrett, E. H., & Loeser, R. 2008, ApJS, 175, 229
  * (10) Baliunas, S. L., Donahue, R. A., Soon, W., & Henry, G. W. 1998, in The Tenth Cambridge Workshop on Cool Stars, Stellar Systems, and the Sun, ASP Conf. Ser. 154, ed. R. A. Donahue & J. A. Bookbinder (San Francisco: ASP), 153
  * (11) Baliunas, S. L., Hartmann, L., & Dupree, A. K. 1983, ApJ, 271, 672
  * (12) Barmby, P., Boyer, M. L., Woodward, C. E., Gehrz, R. D., van Loon J. Th., Fazio, G. G., Marengo, M., & Polomski, E. 2009, AJ, 137, 207
  * (13) Bedin, L. R., Piotto, G., Anderson, J., Cassisi, S., King, I. R., Momany, Y., & Carraro, G. 2004, ApJ, 605, L125
  * (14) Bedin, L. R., Salaris, M., Piotto, G., Cassisi, S., Milone, A. P., Anderson, J., & King, I. R. 2008, ApJ, 679, L29
  * (15) Beers, T. C., Chiba, M., Yoshii, Y., Platais, I., Hanson, R. B., Fuchs, B., & Rossi, S. 2000, AJ, 119, 2866
  * (16) Bianchi, L., Ford, H., Bohlin, R., Paresce, F. & de Marchi, G. 1995, A&A, 301, 537
  * Bond (1980) Bond, H. E. 1980, ApJS, 44, 517
  * (18) Bowen, G. H., & Willson, L. A. 1991, ApJ, 375, L53
  * (19) Boyer, M. L., McDonald, I., van Loon, J. Th., Woodward, C. E., Gehrz, R. D., Evans, A., & Dupree, A. K. 2008, AJ, 135, 1395
  * Boyer et al. (2006) Boyer, M. L., Woodward, C. E., van Loon, J. Th., Gordon, K. D., Evans, A., Gehrz, R. D., Helton, L. A., & Polomski, E. F. 2006, AJ, 132, 1415
  * (21) Buonanno, R., Corsi, C. E., Pecci, F. F., Richer, H. B., & Fahlman, G. G. 1993, AJ, 105, 184
  * (22) Buonanno, R., Corsi, C. E., Pulone, L., Fusi Pecci, F., & Bellazzini, M. 1998, A&A, 333, 505
  * (23) Busso, M., Wasserburg, G. J., Nollett, K. M., & Calandra, A. 2007, ApJ, 671, 802
  * (24) Cacciari, C., et al. 2004, A&A, 413, 343
  * (25) Caloi, V., & D’Antona, F. 2008, ApJ, 673, 847
  * (26) Carney, B. W., Gray, D. F., Yong, D., Latham, D. W., Manset, N., Zelman, R., & Laird, J. B. 2008, AJ, 135, 892
  * (27) Carney, B. W., Latham, D. W., Stefanik, R. P., Laird, J. B., & Morse, J. A. 2003, AJ, 125, 293.
  * (28) Castellani, V., Giannone, P., & Renzini, A. 1971, Ap&SS, 10, 340
  * (29) Catelan, M., Bellazzini, M., Landsman, W. B., Ferraro, F. R., Fusi Pecci, F., & Galleti, S. 2001, AJ, 122, 3171
  * (30) Cenarro, A. J. et al. 2007, MNRAS, 374, 664
  * (31) Cox, D. 2005, ARA&A, 43, 337
  * (32) Cranmer, S. 2008, ApJ, 689, 316
  * (33) D’Antona, F., Caloi, V., Montalbán, J., Ventura, P., & Gratton, R. 2002, A&A, 395, 69
  * (34) Demarque, P., & Eder, J.-A. 1985, in Horizontal-Branch and UV-Bright Stars, ed. A. G. Davis Phillip (Schenectady: Davis), 91
  * (35) Dixon, W. V. D., & Sankrit, R. 2008, ApJ, 686, 1162
  * (36) Dotter, A. 2008, ApJ, 687, L21
  * (37) Dupree, A. K., Avrett, E. H., Brickhouse, N. S., Cranmer, S. R., & Szalai, T. 2008, Cool Stars, Stellar Systems, and the Sun: CS 14 ASP Conf. Ser. 384, ed. G. T. van Belle (San Francisco: ASP), sub77 (CD-ROM) (also astro-ph/0702395)
  * (38) Dupree, A. K., Brickhouse, N. S., Smith, G. H., & Strader, J. 2005, ApJ, 625, L131
  * (39) Dupree, A. K., Hartmann, L., & Avrett, E. H. 1984, ApJ, 281, L37
  * (40) Dupree, A. K., Li, T. Q., & Smith, G. H. 2007, AJ, 134, 1348
  * (41) Dupree, A. K., Sasselov, D. D., & Lester, J. B. 1992a, ApJ, 387, L85
  * (42) Dupree, A. K., & Smith, G. H. 1995, AJ, 110, 405
  * (43) Dupree, A.K., & Whitney, B. A., & Avrett, E. H. 1992b in Cool Stars, Stellar Systems, and the Sun: Proceedings of the Seventh Cambridge Workshop, ASP Conf. Ser. 26, ed. M. S. Giampapa & J. A. Bookbinder (San Francisco: ASP), 525
  * (44) Evans, A., Stickel, M., van Loon, J. Th., Eyres, S. P. S., Hopwood, M. E. L., & Penny, A. J. 2003, A&A, 408, L9
  * (45) Faulkner, D. J., & Freeman, K. C. 1977, ApJ, 211, 77
  * (46) Faulkner, D. J., Scott, T. R., Wood, P. R., & Wright, A. E. 1991, ApJ, 374, L45
  * (47) Frank, J., & Gisler, G. 1976, MNRAS, 176, 533
  * (48) Freire, P. C., Kramer, M., Lyne, A. G., Camilo, F., Manchester, R. N., & D’Amico, N. 2001, ApJ, 557, L105
  * (49) Fulbright, J. P. 2000, AJ, 120, 1841
  * (50) Girardi, L., Bressan, A., Bertelli, G., & Chiosi, C. 2000, A&AS, 141, 371
  * (51) Goldberg, L. 1939, ApJ, 89, 673
  * (52) Gratton, R. G., Sneden, C., Carretta, E., & Bragaglia, A. 2000, A&A, 354, 169
  * (53) Groenewegen, M. A. T., & deJong, T. 1993, A&A, 267, 410
  * (54) Hansen, B. M. S. et al. 2007, ApJ, 671, 380
  * (55) Harris, W. E. 1996, AJ, 112, 1487
  * (56) Hartmann, L. 1998, in Accretion Processes in Star Formation, (New York: Cambridge University Press), Ch. 8
  * (57) Harvey, J. W., & Sheeley, N. R., Jr. 1979, Sp.Sci.Rev., 23, 139
  * (58) Ita, Y., et al. 2007, PASJ, 59, S437
  * (59) Kalirai, J. S., Bergeron, P., Hansen, B. M. S., Kelson, D. D., Reitzel, D. B., Rich, R. M., Richer, H. B. 2007, ApJ, 671, 748
  * (60) Kayser, A., Hilker, M., Grebel, E. K., & Willemsen, P. G., 2008, A&A, 486, 437
  * (61) Koopmann, R. A., Lee, Y.-W., Demarque, P., & Howard, J. M. 1994, ApJ, 423, 380
  * Lambert (1987) Lambert, D. L. 1987, ApJS, 65, 255
  * (63) Lamers, H. J. G. L. M., & Cassinelli, J. P. 1999 in Introduction to Stellar Winds, (New York: Cambridge University Press), Ch. 8
  * (64) Lee, Y.-W., Demarque, P., & Zinn, R. 1994, ApJ, 423, 248
  * (65) Lee, Y.-W., Joo, J.-M., Sohn, Y.-J., Rey, S.-C., Lee, H.-C., & Walker, A. R. 1999, Nature, 402, 55
  * (66) Marshall, J. R., vanLoon, J. Th., Matsuura, M., Wood, P. R., Zijlstra, A. A., & Whitelock, P. A. 2004, MNRAS, 355, 1348.
  * (67) Matsunaga, N. et al. 2008, PASJ, 60, S415
  * (68) Matt, S., & Pudritz, R. E. 2008, ApJ, 678, 1109
  * (69) Mauas, P. J. D., Cacciari, C., & Pasquini, L. 2006, A&A, 454, 609
  * (70) McDonald, I., & van Loon, J. Th. 2007, A&A, 476, 1261
  * (71) McLaughlin, D. E., & van der Marel, R. P. 2005, ApJS, 161, 304
  * (72) McLean, I. S., et al. 1998, SPIE, 3354, 566
  * (73) McLean, I. S., Graham, J. R., Becklin, E. E., Figer, D. F., Larkin, J. E., Levenson, N. A., & Teplitz, H. I. 2000, SPIE, 4008, 1048
  * (74) McLean, I. S., McGovern, M. R., Burgasser, A. J., Kirkpatrick, J. D., Prato, L., & Kim, S. S. 2003, ApJ, 596, 561
  * (75) Mermilliod, J.-C., Mermilliod, M., & Hauck, B. 1997, A&AS, 124, 349
  * (76) Mészáros, Sz., Avrett, E. H., & Dupree, A. K. 2009a, AJ, 138, 615
  * (77) Mészáros, Sz., Dupree, A. K., & Szalai, T. 2009b, AJ, 137, 4282
  * (78) Mészáros, Sz., Dupree, A. K., & Szentgyorgyi, A. 2008, AJ, 135, 1117
  * (79) Moehler, S., Koester, D., Zoccali, M., Ferraro, F. R., Heber, U., Napiwotzki, R., & Renzini, A. 2004, A&A, 420, 515
  * (80) Nordhaus, J., Busso, M., Wasserburg, G. J., Blackman, E. G., & Palmerini, S. 2008, ApJ, 684, L29
  * O’Brien & Lambert (1986) O’Brien, G. T., Jr., & Lambert, D. L. 1986, ApJS, 62, 899
  * (82) Okada, Y., Kokubun, M., Yuasa, T., & Makishima, K. 2007, PASJ, 59, 727
  * Origlia et al. (2002) Origlia, L., Ferraro, F. R., Fusi Pecci, F., & Rood, R. T. 2002, ApJ, 571, 458
  * Origlia et al. (2007) Origlia, L., Rood, R. T., Fabbri, S., Ferraro, F. R., Fusi Pecci, F., & Rich, R. M. 2007, ApJ, 667, L85
  * (85) Pancino, E., Ferraro, F. R., Bellazzini, M., Piotto, G., & Zoccali, M. 2000, ApJ, 534, L83
  * (86) Peterson, R. C., Rood, R. T., & Crocker, D. A. 1995, ApJ, 453, 214
  * (87) Pilachowski, C. A., Sneden, C., & Booth, J. 1993, ApJ, 407, 699
  * Pilachowski et al. (1996) Pilachowski, C. A., Sneden, C., & Kraft, R. P. 1996, AJ, 111, 1689
  * (89) Piotto, G., et al. 2007, ApJ, 661, L53
  * (90) Preston, G. W., 1997, AJ, 113, 1860
  * (91) Recio-Blanco, A., Aparicio, A., Piotto, G., de Angeli, F., & Djorgovski, S. G. 2006, A&A, 452, 875
  * (92) Reimers, D. 1975, Mem. Soc. R. Sci. Liege, Ser. 6, 8, 369
  * (93) Renzini, A. 1981, in Phys. Proc. in Red Giants, ed. I. Iben & A. Renzini, (Dordrecht: Reidel), 431
  * (94) Richardson, J. D., Kasper, J. C., Wang, C., Belcher, J. W., & Lazarus, A. J. 2008, Nature, 454, 63
  * (95) Richer, H. B., et al. 2008, AJ, 135, 2141
  * (96) Rood, R. T. 1973, ApJ, 184, 815
  * (97) Rossi, S., Beers, T. C., Sneden, C., Sevastyanenko, T., Rhee, J., & Marsteller, B. 2005, AJ, 130, 2804
  * (98) Sandage, A., & Wildey, R. 1967, ApJ, 150, 469
  * (99) Sandquist, E. L., & Martel, A. R. 2007, ApJ, 654, L65
  * (100) Sanz-Forcada, J., & Dupree, A. K. 2008, A&A, 488, 715
  * (101) Sasselov, D. D., & Lester, J. B. 1994, ApJ, 423, 795
  * (102) Schwarzschild, M. 1970, Quart, J. R. A. S. 11, 12
  * (103) Sills, A., & Pinsonneault, M. H. 2000, ApJ, 540, 489
  * (104) Smith, G. H. 1998, PASP, 110, 1119
  * (105) Smith, G. H. 1999, PASP, 111, 980
  * (106) Smith, G. H., & Dupree, A. K. 1988, AJ, 95, 1547
  * (107) Smith, G. H., Dupree, A. K., & Strader, J. 2004, PASP, 116, 819
  * (108) Sneden, C., Kraft, R. P., Guhathakurta, P., Peterson, R. C., & Fulbright, J. P. 2004, AJ, 127, 2162
  * (109) Soker, N., Catelan, M., Rood, R. T., & Harpaz, A. 2001a, ApJ, 563, L69
  * (110) Soker, N., Rappaport, S., & Fregeau, J. 2001b, ApJ, 563, L87
  * (111) Sweigart, A. V. 1997, ApJ, 474, L23
  * (112) Tayler, R. J., & Wood, P. R. 1975, MNRAS, 171, 467
  * (113) Tokunaga, A. T. 2000, in Allen’s Astrophysical Quantities, 4th ed., ed. A. N. Cox (New York: Springer), 143
  * (114) VandenBerg, D. A., & Faulkner, D. J. 1977, ApJ, 218, 415
  * (115) VandenBerg, D. A., Bergbusch, P. A., & Dowler, P. D. 2006, ApJS, 162, 375
  * (116) van Loon, J. Th., Boyer, M. L., & McDonald, I., 2008, ApJ, 680, L49
  * (117) van Loon, J. Th., Stanimirović, S., Evans, A., & Muller, E. 2006, MNRAS, 365, 1277
  * (118) Ventura, P., & D’Antona, F. 2005, A&A, 439, 1075
  * (119) Vink, J. S., & Cassisi, S. 2002, A&A, 392, 553
  * (120) Wang, Y.-M. 1998, in Cool Stars, Stellar Systems, and the Sun: Tenth Cambridge Workshop, ed. R. A. Donahue & J. A. Bookbinder, ASP Conf. Ser. 154, 131
  * (121) Yong, H., Demarque, P., & Yi, S. 2000, ApJ, 539, 928
  * (122) Zarro, D. M., & Zirin, H. 1986, ApJ, 304, 365
  * (123) Zirin, H. 1975, ApJ, 199, L63
  * (124) Zirin, H. 1982, ApJ, 260, 655

Figure 1: Red giant (RGB) stars showing the region of the He I line. The Si I
photospheric line is the strong narrow absorption at 10827.09Å. The position
and extent of the He I triplet line spanning 10830.081-10830.341Å are marked.
The strongest member of the multiplet occurs at 10830.341Å. Weak narrow
emission features sometimes arise from the incomplete sky subtraction. A bad
pixel causes the narrow emission spike in BD+09°2870.

Figure 2: Red horizontal branch stars in the sample. See caption for Figure 1

Figure 3: The He I line in normalized spectra of the 6 AGB stars in our
sample. The position of the He I $\lambda$10830.341 transition is marked by a
broken line set to zero velocity. The continuum level is set at 1.0 for each
star, with each spectrum offset by a constant value; the extent of 0.1 in the
continuum level is shown. The strong Si I photospheric line at $-$90 km s-1
dominates this region of the spectrum. The AGB stars are arranged in order
(lower to upper spectra) of increasing $(B-V)_{0}$ values which are noted
below the stellar identification. The most luminous objects show weak emission
in He I. Evidence of absorption at high negative velocities is seen in the 4
coolest stars: BD+12°2547, HD 121135, HD 107752, BD+18°2757.

Figure 4: He I spectra for the two subgiants in the sample.

Figure 5: HD 135148, a CH star, shows a substantial deep P Cygni profile with
extent at least to the Si I line and beyond that indicates a wind velocity of
$\sim$ 115 km s-1.

Figure 6: The He spectral region of HD 104207 (GK Com), the coolest red giant
in this sample, where the spectrum is dominated by neutral lines of Si, Ti,
Fe, and Ca. In addition, the Ti I line at 10828.04Å occurs in very cool stars
near the Si I transition at 10827.09Å. The He I line at $\lambda$10830.34
appears to have weak emission longward of the absorption. The short wavelength
wing of the Si I 10827.09Å profile may have additional absorption when
compared to the other Si I transition at 10843.9Å, possibly caused by extended
helium absorption.

Figure 7: Location of target stars in a color-magnitude diagram (top panel)
and in a Teff-magnitude diagram (lower panel). Results for 6 giants in M13 and
5 metal-poor field giants reported previously (Dupree et al. 1992a; Smith et
al. 2004) have been added to this figure (see Table 3). The curves mark 12.1
Gyr isochrones (VandenBerg et al. 2006) for an abundance [Fe/H]=$-$2.01 (solid
line) and [Fe/H]=$-$1.53 (broken line). Plus signs mark both P Cygni profiles
and asymmetric profiles signaling mass outflow.

Figure 8: Equivalent width of the He I absorption line as a function of
$T_{eff}$ with evolutionary status of the stars indicated as follows: RGB=red
giant branch; AGB= asymptotic giant branch; SGB= subgiant branch. The plus
($+$) signs overplotted indicate stars with P Cygni profiles. Lower limits to
the equivalent widths occur when the helium absorption overlaps the Si I
photospheric line.

Figure 9: Equivalent width of $\lambda$10830 in the RHB sample as a function
of $T_{eff}$. Absorption becomes systematically stronger in stars with
$T_{eff}\gtrsim$ 5320K. The plus ($+$) sign overplotted indicates a star with
a P Cygni profile. Lower limits to the equivalent widths occur when the helium
absorption overlaps the Si I photospheric line. Errors in the equivalent width
amount to $\sim\pm$5%.

Figure 10: Equivalent widths of the He I absorption in red giants and
subgiants from our sample and from the red giants in Population I stars
reported by O’Brien & Lambert (1987). Many of the Population I objects were
observed several times and the lower and upper values of the equivalent widths
are connected by solid lines. The plus ($+$) sign overplotted indicates a P
Cygni profile. Some Population I giants with strong X-ray emission have
substantially larger equivalent widths (O’Brien & Lambert 1986; Sanz-Forcada &
Dupree 2008), and these are omitted from this figure.

Figure 11: Equivalent widths of He I absorption as a function of [Fe/H] with
evolutionary status of the stars indicated, as follows: RGB=red giant branch;
RHB=red horizontal branch; AGB= asymptotic giant branch; SGB= subgiant branch.
The CH star, HD 135148 with EW= 2390 mÅ and [Fe/H]=$-$1.9 has been omitted as
well as 3 stars (2 RGB and TY Vir) not showing helium. Stars displaying a P
Cygni profile are marked with a plus ($+$) sign; lower limits are shown where
the helium absorption extends into the neighboring Si I line. There is no
systematic dependence of the equivalent width on [Fe/H] between [Fe/H]=$-$0.7
and $-$3.0. Figure 12: The ratio of the short wavelength velocity extent
(BLUE) to the long wavelength velocity extent (RED) of the absorption profile
of the He I 10830Å line as a function of absolute visual magnitude. Only stars
with helium absorption lines and no emission are included here. A line arising
in a stationary atmosphere has a ratio of 1; outflow is indicated when
Blue/Red $>$1\. An ’up’ arrow marks stars for which the blue velocity wing
reached the Si I line, 90 km s-1 to shorter wavelengths, but the limiting
extent could not be determined because of overlap with the Si I feature. The
error bar marks the estimated 15% uncertainty in the measurement of the ratio.
The majority of these stars have asymmetric absorption profiles signalling
outflow.

Figure 13: Short wavelength extent of the helium 10830Å line as a function of
[Fe/H]. An extent of $\sim$40 km s-1 appears to correspond to the normal
thermal and turbulent width of the line (marked by the broken line). No
systematic dependence is present as a function of metallicity. Stars
displaying a P Cygni profile are marked with a plus ($+$) sign. Upward
pointing arrows denote lower limits because the helium absorption extends into
the neighboring Si I line and the extent of the helium profile can not be
determined.

Figure 14: The relation between the bolometric absolute magnitude, $M_{bol}$,
and the mass loss rate (Equation 7) as inferred from the He I 10830Å profile
in the stars with $V_{term}$ exceeding the thermal width of 45 km s-1. Upper
and lower limits are shown that derive from the limits on the the lower level
population ratio. Some stars exhibit absorption extending into the Si I line
located $-$90 km s-1 from He I. Since the termination of the He I profile is
difficult to establish, the figure shows a lower limit arrow on these rates.
Figure 15: The depth of the He I 10830Å absorption as a function of the ratio
of the lower level population, n(${}^{3}S$), in He I to the total helium
abundance, n(Hetotal), as calculated for various stellar chromospheres. Here,
Hetotal= He I + He II. The vertical bars on two stars indicate the span of
line-core depth values at different $\mu$ where $\mu$=cos $\theta$, the polar
angle with respect to the stellar photosphere in the plane-parallel
approximation that was used. With the exception also of the Sun, all of the
other models were calculated in a spherical approximation. The giant star
model has parameters: $T_{eff}$ = 4250K, log g = 1.6, and solar abundances
other than hydrogen and helium are decreased by a factor of 0.33. From these
results, we take an upper limit of the lower level population ratio,
n(${}^{3}S$)/n($He_{total}$)= $-$6.2 dex, and a lower limit of this ratio of
$-$7.0 dex.  Table 1: Metal-Poor Field Giants Observed

Star | $V$ | $B-V$ | $U-B$ | Ref. | $E(b-y)$ | $M_{V}$ | $V$ | $b-y$ | $c_{1}$ | Ref. | E($B-V$) | ($B-V$)0 | [Fe/H] | Ref. | Evol.aaEvolutionary state: RGB = red giant branch; SGB = subgiant branch; SR= semiregular variable; AGB = asymptotic giant branch decided on the basis of (b$-$y, c1) diagram and $M_{V}$; RHB = red horizontal branch decided on the basis of (b$-$y, c1) diagram and $M_{V}$.
---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---
BD +01 3070 | 10.060 | 0.769 |  $\cdots$ | 3 | 0.016 | 1.46 | 10.038 | 0.487 | 0.264 | 2 | 0.022 | 0.747 |  $-$1.85 | 2 | RGB
BD +05 3098 | 10.530 | 0.780 |  $\cdots$ | 3 | 0.028 |  $-$0.04 | 10.537 | 0.542 | 0.380 | 2 | 0.038 | 0.742 |  $-$2.4 | 2 | RGB
BD +09 2574 | 10.517 | 0.793 |  $\cdots$ | 3 | 0.000 | 0.27 | 10.523 | 0.530 | 0.379 | 2 | 0.000 | 0.793 |  $-$1.95 | 2 | RGB
BD +09 2860 | 10.830 | 0.710 |  $\cdots$ | 3 | 0.003 | 0.73 | 10.83 | 0.440 | 0.430 | 2 | 0.004 | 0.706 |  $-$1.6 | 2 | RHB
BD +09 2870 | 9.440 | 1.042 |  $\cdots$ | 3 | 0.012 |  $-$1.27 | 9.426 | 0.647 | 0.527 | 2 | 0.016 | 1.026 |  $-$2.37 | 6 | RGB
BD +09 3223 | 9.260 | 0.670 |  $\cdots$ | 3 | 0.041 | 0.58 | 9.253 | 0.469 | 0.481 | 2 | 0.056 | 0.614 |  $-$2.26 | 6 | RHB
BD +10 2495 | 9.723 | 0.749 |  $\cdots$ | 3 | 0.002 | 0.29 | 9.745 | 0.522 | 0.361 | 2 | 0.003 | 0.746 |  $-$1.83 | 7 | RGB
BD +11 2998 | 9.067 | 0.679 |  $\cdots$ | 3 | 0.035 | 0.77 | 9.058 | 0.453 | 0.505 | 2 | 0.048 | 0.631 |  $-$1.17 | 6 | RHB
BD +12 2547 | 9.920 | 1.034 |  $\cdots$ | 3 | 0.004 |  $-$0.94 | 9.92 | 0.635 | 0.420 | 2 | 0.005 | 1.029 |  $-$0.72 | 7 | AGB
BD +17 3248 | 9.37 | 0.66 | 0.08 | 1 | 0.040 | 0.65 | 9.352 | 0.486 | 0.445 | 2 | 0.055 | 0.605 |  $-$2.02 | 6 | RHB
BD +18 2757 | 9.795 | 0.745 | 0.155 | 1 | 0.000 |  $-$0.80 | 9.84 | 0.550 | 0.490 | 2 | 0.00 | 0.745 |  $-$2.19 | 6 | AGB
BD +18 2976 | 9.850 | 1.051 |  $\cdots$ | 3 | 0.005 |  $-$1.32 | 9.835 | 0.655 | 0.527 | 2 | 0.006 | 1.045 |  $-$2.4 | 2 | RGB
BD +30 2611 | 9.125 | 1.240 | 1.125 | 1 | 0.003 |  $-$1.11 | 9.143 | 0.807 | 0.551 | 2 | 0.004 | 1.236 |  $-$1.49 | 6 | RGB
BD +52 1601 | 8.800 | 0.901 |  $\cdots$ | 3 | 0.000 | 0.13 | 8.80 | 0.555 | 0.445 | 2 | 0.00 | 0.901 |  $-$1.58 | 6 | RGB
BD +54 1323 | 9.343 | 0.670 |  $\cdots$ | 3 | 0.003 | 0.69 | 9.33 | 0.470 | 0.440 | 2 | 0.004 | 0.666 |  $-$1.65 | 6 | RHB
BD +58 1218 | 9.960 | 0.835 |  $\cdots$ | 3 | 0.000 | 0.39 | 9.96 | 0.515 | 0.360 | 2 | 0.00 | 0.835 |  $-$2.66 | 7 | RGB
BD $-$03 5215 | 10.170 | 0.719 |  $\cdots$ | 3 | 0.032 | 0.73 | 10.188 | 0.435 | 0.499 | 2 | 0.043 | 0.676 |  $-$1.66 | 1 | RHB
CD $-$30 8626 | 9.703 | 0.759 | 0.250 | 1 | 0.034 | 0.27 | 9.719 | 0.521 | 0.479 | 2 | 0.047 | 0.712 |  $-$1.7 | 2 | AGB
HD 83212 | 8.335 | 1.070 | 0.760 | 1 | 0.018 |  $-$0.92 | 8.328 | 0.694 | 0.571 | 2 | 0.025 | 1.045 |  $-$1.49 | 2 | RGB
HD 93529 | 9.300 | 0.881 | 0.370 | 1 | 0.048 | 1.05 | 9.306 | 0.582 | 0.409 | 2 | 0.066 | 0.815 |  $-$1.67 | 6 | RGB
HD 101063 | 9.460 | 0.755 | 0.098 | 1 | 0.027 | 2.74 | 9.47 | 0.499 | 0.284 | 2 | 0.037 | 0.718 |  $-$1.13 | 7 | SGB
HD 104207 | 6.984 | 1.574 | 1.565 | 1 |  $\cdots$ |  $-$2.48 |  $\cdots$ |  $\cdots$ |  $\cdots$ | 4 | 0.04 | 1.534 |  $-$1.93 | 4 | RGB
HD 105546 | 8.616 | 0.708 |  $\cdots$ | 3 | 0.000 | 0.79 | 8.61 | 0.460 | 0.420 | 2 | 0.000 | 0.708 |  $-$1.27 | 6 | RHB
HD 107752 | 10.05 | 0.75 | 0.18 | 1 | 0.000 |  $-$0.68 | 9.994 | 0.578 | 0.463 | 2 | 0.000 | 0.75 |  $-$2.88 | 7 | AGB
HD 108317 | 8.038 | 0.631 |  $\cdots$ | 3 | 0.000 | 0.52 | 8.044 | 0.450 | 0.311 | 2 | 0.000 | 0.631 |  $-$2.24 | 6 | RHB
HD 108577 | 9.597 | 0.694 | 0.134 | 1 | 0.015 |  $-$0.57 | 9.581 | 0.506 | 0.500 | 2 | 0.021 | 0.673 |  $-$2.28 | 6 | AGB
HD 110184 | 8.305 | 1.175 | 0.765 | 1 | 0.000 |  $-$2.14 | 8.293 | 0.818 | 0.712 | 2 | 0.000 | 1.175 |  $-$2.56 | 6 | RGB
HD 110885 | 9.180 | 0.672 |  $\cdots$ | 3 | 0.000 | 0.74 | 9.18 | 0.423 | 0.492 | 2 | 0.000 | 0.672 |  $-$1.44 | 8 | RHB
HD 111721 | 7.971 | 0.799 | 0.157 | 1 | 0.006 | 1.16 | 7.98 | 0.526 | 0.315 | 2 | 0.008 | 0.791 |  $-$1.26 | 8 | RGB
HD 113002 | 8.745 | 0.747 | 0.209 | 1 |  $\cdots$ | 2.95 |  $\cdots$ |  $\cdots$ |  $\cdots$ | 5 | 0.020 | 0.727 |  $-$1.08 | 5 | SGB
HD 115444 | 8.967 | 0.784 | 0.173 | 1 | 0.000 |  $-$0.49 | 8.98 | 0.575 | 0.425 | 2 | 0.000 | 0.784 |  $-$2.77 | 6 | RGB
HD 119516 | 9.090 | 0.661 |  $\cdots$ | 3 | 0.000 | 0.56 | 9.09 | 0.410 | 0.525 | 2 | 0.000 | 0.661 |  $-$2.5 | 2 | RHB
HD 121135 | 9.357 | 0.795 |  $\cdots$ | 3 | 0.008 |  $-$0.36 | 9.368 | 0.530 | 0.509 | 2 | 0.011 | 0.784 |  $-$1.57 | 6 | AGB
HD 122956 | 7.22 | 1.01 |  $\cdots$ | 1 | 0.042 |  $-$0.69 | 7.251 | 0.667 | 0.480 | 2 | 0.058 | 0.95 |  $-$1.78 | 6 | RGB
HD 126587 | 9.125 | 0.818 | 0.160 | 1 | 0.058 |  $-$0.11 | 9.097 | 0.596 | 0.381 | 2 | 0.079 | 0.739 |  $-$3.06 | 7 | RGB
HD 126778 | 8.168 | 0.916 | 0.666 | 1 | 0.000 | 1.05 | 8.15 | 0.596 | 0.449 | 2 | 0.000 | 0.916 |  $-$0.7 | 2 | RGB
HD 135148 | 9.490 | 1.388 |  $\cdots$ | 3 | 0.016 |  $-$0.87 | 9.425 | 0.869 | 0.440 | 2 | 0.022 | 1.366 |  $-$1.90 | 6 | RGBbbCH star.
HD 141531 | 9.130 | 1.240 |  $\cdots$ | 3 | 0.012 |  $-$1.46 | 9.145 | 0.765 | 0.603 | 2 | 0.016 | 1.224 |  $-$1.62 | 8 | RGB
HD 161770 | 9.681 | 0.665 |  $-$0.041 | 1 | 0.054 | 1.45 | 9.70 | 0.500 | 0.281 | 2 | 0.074 | 0.591 |  $-$2.12 | 7 | RGB
HD 195636 | 9.540 | 0.645 |  $-$0.005 | 1 | 0.044 | 0.51 | 9.552 | 0.467 | 0.481 | 2 | 0.060 | 0.585 |  $-$2.83 | 10 | RHB
TY Vir | 8.100 | 1.28 | 1.00 | 1 | 0.012 |  $-$1.17 | 8.165 | 0.938 | 0.711 | 2 | 0.016 | 1.26 |  $-$1.78 | 9 | SR

References. — 1. Mermilliod et al. 1997; 2. Anthony-Twarog & Twarog 1994; 3.
HIPPARCOS Input Catalogue; 4. GK Com (Var.); M4 III; Beers et al. 2000 data
used to derive $M_{V}$; 5. Beers et al. (2000) data used to derive $M_{V}$; 6.
Pilachowski et al. (1996); 7. Pilachowski et al. (1993); 8\. Gratton et al.
2000; 9. Fulbright 2000; 10. Anthony-Twarog & Twarog 1998.

Table 2: Parameters of He I 10830Å Line

Star | Exp.aaTotal exposure time, usually divided into several nodded segments. | Evol.bbEvolutionary stage estimated from photometry: RGB= red giant branch; AGB=asymptotic giant branch; RHB= red horizontal branch; SGB= sub-giant branch; SR= semi-regular variable. | RV | Teff | LineccCode for the presence of He I $\lambda$10830; 0: no He line observed; 1: He absorption; 2: P Cygni profile or emission observed (2 stars). | EWddEquivalent Width: Positive values indicate equivalent width of absorption line below the local continuum; negative values indicate equivalent width of the emission line. | VtermeeFurthest extent of the short wavelength absorption edge. |  $B/R$ff$B/R$ is the ratio of the short wavelength extent of helium absorption ($V_{term}$) to the long wavelength extent.  | R⋆ | Vesc(2R⋆)ggStellar mass assumed: RGB=0.75M⊙; RHB = 0.7M⊙; AGB =0.6M⊙; SGB=0.8M⊙; and SR=0.6M⊙. | Ref.
---|---|---|---|---|---|---|---|---|---|---|---
| (s) |  | (km s-1) | (K) |  | (mÅ) | (km s-1) |  | (R⊙) | (km s-1) | (for Teff)
BD +01 3070 | 400 | RGB |  $-$329.9 | 5130 | 1 | 146.3 |  $\gtrsim$90 |  $\gtrsim$0.80 | 6.4 | 150 | 1
BD +05 3098 | 1440 | RGB |  $-$160.5 | 4930 | 1 | 51.5 | 51 | 1.19 | 14.3 | 100 | 1
BD +09 2574 | 1200 | RGB |  $-$49.8 | 4860 | 2 | 56.1 | 52 | 0.55 | 12.8 | 106 | 1
BD +09 2860 | 2160 | RHB |  $-$20.7 | 5240 | 1 | 38.2 | 44 | 0.97 | 8.5 | 126 | 1
BD +09 2870 | 360 | RGB |  $-$120.1 | 4600 | 2 | 15.9 | 60 | 0.84 | 30.7 | 69 | 2
BD +09 3223 | 720 | RHB | 67.3 | 5310 | 1 | 25.5 | 44 | 1.08 | 8.9 | 123 | 1
BD +10 2495 | 840 | RGB | 263.2 | 4920 | 1 | 25.2 | 46 | 1.11 | 12.3 | 108 | 1
BD +11 2998 | 480 | RHB | 50.7 | 5360 | 1 | 54.1 | 45 | 0.47 | 7.8 | 131 | 1
BD +12 2547 | 420 | AGB | 5.3 | 4610 | 2 | 8.8: |  $\gtrsim$90 |  $\gtrsim$0.86 | 26.0 | 67 | 1
BD +17 3248 | 720 | RHB |  $-$147.4 | 5250 | 1 | 45.4 | 62 | 1.25 | 8.8 | 123 | 2
BD +18 2757 | 500 | AGB |  $-$29.0 | 4840 | 1 | 86.3 |  $\gtrsim$90 |  $\gtrsim$2.76 | 21.3 | 74 | 1
BD +18 2976 | 500 | RGB |  $-$167.4 | 4550 | 2 | 17.2: | 143 | 2.65 | 32.5 | 67 | 1
BD +30 2611 | 200 | RGB |  $-$282.8 | 4275 | 1 | 92.6 | 143 | 1.78 | 35.9 | 63 | 2
BD +52 1601 | 400 | RGB |  $-$47.4 | 4750 | 0 | 0. |  $\cdots$ |  $\cdots$ | 14.6 | 99 | 2
BD +54 1323 | 600 | RHB |  $-$67.2 | 5300 | 1 | 45.0 | 36 | 0.48 | 8.4 | 126 | 2
BD +58 1218 | 600 | RGB |  $-$305.2 | 4950 | 1 | 37. | 65 | 1.54 | 11.7 | 111 | 3
BD $-$03 5215 | 1200 | RHB |  $-$294.5 | 5420 | 1 | 232.4 |  $\gtrsim$90 |  $\gtrsim$1.39 | 7.8 | 131 | 1
CD $-$30 8626 | 600 | AGB | 266.2 | 5000 | 1 | 25.6 | 41 | 1.08 | 11.9 | 99 | 1
HD 83212 | 500 | RGB | 108.0 | 4550 | 2hhEmission clearly present in HD 141531 but absorption not apparent. |  $-$43.6 |  $\gtrsim$90 |  $\gtrsim$1.42 | 26.9 | 73 | 2
HD 93529 | 300 | RGB | 145.6 | 4650 | 1 | 39.9 | 38 | 0.79 | 10.2 | 119 | 2
HD 101063 | 400 | SGB | 182.6 | 5070 | 1 | 68.3 | 48 | 1.04 | 3.6 | 205 | 3
HD 104207 | 6 | RGB | 35.6 | 3916 | 2 | 23.8 |  $\gtrsim$90 |  $\gtrsim$8.69 | 95.0 | 39 | 4
HD 105546 | 240 | RHB | 18.1 | 5300 | 1 | 31.7 | 35 | 0.99 | 8.0 | 130 | 2
HD 107752 | 600 | AGB | 219.2 | 4750 | 1 | 71.7 |  $\gtrsim$90 |  $\gtrsim$0.92 | 21.4 | 73 | 1
HD 108317 | 150 | RHB | 4.4 | 5230 | 1 | 23.2 | 35 | 1.11 | 9.5 | 119 | 1
HD 108577 | 400 | AGB |  $-$112.1 | 4975 | 1 | 20.6 | 42 | 1.11 | 17.8 | 80 | 2
HD 110184 | 120 | RGB | 138.5 | 4250 | 0 | 0. |  $\cdots$ |  $\cdots$ | 59.0 | 49 | 2
HD 110885 | 600 | RHB |  $-$48.8 | 5330 | 1 | 14.1 |  $\gtrsim$90 |  $\gtrsim$0.85 | 8.1 | 129 | 1
HD 111721 | 240 | RGB | 20.5 | 5080 | 1 | 48.4 | 50 | 1.11 | 7.5 | 138 | 5
HD 113002 | 240 | SGB |  $-$95.2 | 5007 | 1 | 30.1 | 57 | 1.46 | 3.4 | 212 | 4
HD 115444 | 300 | RGB |  $-$27.7 | 4750 | 2 | 32.6 | 65 | 0.71 | 19.6 | 86 | 2
HD 119516 | 600 | RHB |  $-$287.2 | 5440 | 1 | 466. | 121 | 1.29 | 8.6 | 125 | 1
HD 121135 | 600 | AGB | 125.0 | 4925 | 2 | 63.2 | 104 | 1.47 | 16.5 | 84 | 2
HD 122956 | 70 | RGB | 165.2 | 4600 | 1 | 84. | 110 | 2.99 | 23.4 | 78 | 2
HD 126587 | 480 | RGB | 149.3 | 4960 | 1 | 20.9 | 39 | 0.97 | 14.7 | 99 | 5
HD 126778 | 150 | RGB |  $-$138.8 | 4847 | 2 | 33.8 | 41 | 0.56 | 8.9 | 127 | 6
HD 135148 | 320 | RGB |  $-$85.4 | 4275 | 2 | 2389.6 | 117 | 1.51 | 32.2 | 67 | 2
HD 141531 | 440 | RGB | 2.3 | 4340 | 2ggStellar mass assumed: RGB=0.75M⊙; RHB = 0.7M⊙; AGB =0.6M⊙; SGB=0.8M⊙; and SR=0.6M⊙. |  $-$12.8 |  $\cdots$ |  $\cdots$ | 40.1 | 60 | 1
HD 161770 | 1020 | RGB |  $-$130.6 | 5406 | 1 | 30.5 | 32 | 0.91 | 5.7 | 159 | 4
HD 195636 | 600 | RHB |  $-$257.7 | 5370 | 2 | 313.7 | 171 | 1.76 | 9.1 | 122 | 1
TY Vir | 70 | SR | 229.1 | 4350 | 0 | 0. |  $\cdots$ |  $\cdots$ | 34.8 | 58 | 8

References. — 1. Carney et al. 2003; 2. Pilachowski et al. 1996; 3\. Carney et
al. 2008; 4. Alonso et al. 1999; 5. Rossi et al. 2005; 6\. Cenarro et al. 2007
8. Andrievsky et al. 2007

Table 3: He I 10830Å Observations from Previous Publications

Star | MV | (B$-$V)0 | [Fe/H] | He 10830Å | VtermaaFurthest extent of the short wavelength absorption edge. | Teff | R⋆ | Vesc(2R⋆)bbStellar mass assumed: RGB=0.75M⊙; RHB=0.7M⊙; AGB =0.6M⊙. | Ref.
---|---|---|---|---|---|---|---|---|---
|  |  |  |  | (km s-1) | (K) | (R⊙) | (km s-1) |
Field Giants
BD +17°3248ccRHB star. | +0.65 | 0.605 |  $-$2.1 | absorption | 60 | 4625 | 8.8 | 123 | 2
HD 6833 |  $-$0.9 | 1.08 |  $-$0.91 | absorption |  $\gtrsim$90 | 4400 | 29.6 | 70 | 1,3
HD 122563 |  $-$1.24 | 0.90 |  $-$2.6 | absorption | 140 | 4625 | 29.6 | 70 | 2
HD 165195 |  $-$2.14 | 1.07 |  $-$2.4 | not detect. |  $\cdots$ | 4450 |  $\cdots$ |  $\cdots$ | 2
HD 221170 |  $-$1.67 | 1.09 |  $-$2.0 | not detect. |  $\cdots$ | 4425 |  $\cdots$ |  $\cdots$ | 2
Red Giants in M13
II-33 |  $-$1.78 | 1.20 |  $-$1.51 | not detect. |  $\cdots$ | 4390 |  $\cdots$ |  $\cdots$ | 2
III-37 |  $-$1.67 | 1.14 |  $-$1.51 | not detect. |  $\cdots$ | 4400 |  $\cdots$ |  $\cdots$ | 2
III-63 |  $-$2.25 | 1.37 |  $-$1.51 | not detect. |  $\cdots$ | 4200 |  $\cdots$ |  $\cdots$ | 2
III-73 |  $-$2.13 | 1.28 |  $-$1.51 | not detect. |  $\cdots$ | 4300 |  $\cdots$ |  $\cdots$ | 2
IV-15ddAGB star. |  $-$1.49 | 1.02 |  $-$1.51 | absorption | 30 | 4650 | 32.7 | 59 | 2
IV-25 |  $-$2.36 | 1.52 |  $-$1.51 | not detect. |  $\cdots$ | 4000 |  $\cdots$ |  $\cdots$ | 2

References. — 1. Dupree et al. 1992a; 2. Smith et al. 2004 3. Smith & Dupree
1988