# Direct Evaluation of the Helium Abundances in Omega Centauri

A. K. Dupree and E. H. Avrett  Harvard-Smithsonian Center for Astrophysics,
Cambridge, MA 02138, USA<EMAIL_ADDRESS><EMAIL_ADDRESS>
###### Abstract

A direct measure of the helium abundances from the near-infrared transition of
He I at 1.08$\mu$m is obtained for two nearly identical red giant stars in the
globular cluster Omega Centauri (catalog ). One star exhibits the He I line;
the line is weak or absernt in the other star. Detailed non-LTE semi-empirical
models including expansion in spherical geometry are developed to match the
chromospheric H$\alpha$, H$\beta$, and Ca II K lines, in order to predict the
helium profile and derive a helium abundance. The red giant spectra suggest a
helium abundance of $Y\leq 0.22$ (LEID 54064) and $Y=0.39-0.44$ (LEID 54084)
corresponding to a difference in the abundance $\Delta Y\geq 0.17$. Helium is
enhanced in the giant star (LEID 54084) that also contains enhanced aluminum
and magnesium. This direct evaluation of the helium abundances gives
observational support to the theoretical conjecture that multiple populations
harbor enhanced helium in addition to light elements that are products of
high-temperature hydrogen burning. We demonstrate that the 1.08$\mu$m He I
line can yield a helium abundance in cool stars when constraints on the semi-
empirical chromospheric model are provided by other spectroscopic features.

stars: individual (LEID 54064 (catalog Cl*NGC5139 LEID54064), LEID 54084
(catalog Cl*NGC5139 LEID54084) ) - stars: abundances - stars: atmospheres -
globular clusters: individual (Omega Centauri)

††slugcomment: ApJ Letters, in press

##  1 Introduction

Our current understanding of stellar populations in globular clusters has
dramatically changed with the discoveries of multiple stellar generations in a
single globular cluster. While variations in color and a spread in the [Fe/H]
values of red giants in massive clusters have been long recognized (Woolley
1966, Geyer 1967) along with variations of light elements (Martell 2011), the
firm identification of multiple populations on the main sequence in Omega
Centauri (Anderson 1997; Bedin et al. 2004; Bellini et al. 2010), and
subsequently several other clusters (cf. Gratton et al. 2012), was surprising
and continues to present theoretical challenges. Norris (2004) suggested,
based on isochrone calculations, that dwarf stars on the ‘blue’ main sequence
in Omega Cen would be enhanced in helium by $\Delta$Y$\sim$0.10$-$0.15. The
lowered hydrogen opacity causes stars of the same mass to appear hotter and
more luminous (Valcarce et al. 2012). Subsequently, the assessment of metals
in dwarfs on the bifurcated main sequence in Omega Cen, showed that the hotter
objects (the ‘blue’ dwarfs) were less metal-poor than the ‘red’ dwarf stars
(Piotto et al. 2005). Stellar models suggest that increased metals also signal
the presence of enhanced helium in the ‘blue’ main sequence. The source (or
sources) of such an enhancement remains elusive. One attractive explanation
appears to be a second stellar generation formed from the material lost by the
first generation of intermediate mass stars during their asymptotic giant
phases (D’Ercole et al. 2010; Johnson & Pilachowski 2010; Renzini 2013),
although other possibilities such as fast-rotating massive stars (Charbonnel
et al. 2013) or massive binary-star mass overflow (de Mink et al. 2009) may
well contribute (cf. Gratton et al. 2012). The formation of cluster
populations with several generations of star formation also impacts an
understanding of the halo of the Milky Way, satellites of our Galaxy, and the
star formation and assembly history of other galaxies (cf. Gratton et al.
2012; Brodie & Strader 2006).

It is obviously of great interest if a helium enhancement could be verified in
globular cluster stars in our Galaxy. A direct measure of the helium abundance
from a spectrum would provide confirmation of Norris’ conjecture. Such a
measurement is challenging because useful lines of helium are generally absent
in the optical spectra of cool stars. Moreover, in hotter stars, such as blue
horizontal branch objects, sedimentation caused by diffusion and element
stratification occur. Helium abundances from the spectroscopy of hot
horizontal branch stars in Omega Cen demonstrate the effects of surface
diffusion, or mixing during late helium core flashes (DaCosta et al. 1986;
Moehler et al. 2011; Moni Bidin et al. 2012) and derived abundance values vary
widely from Y $\leq$0.02 to Y=0.9.

In cool stars, a transition in He I occurs in the near-infrared at 1.08$\mu$m
and has been identified in many metal-poor field stars, where, in addition to
abundances, it can indicate atmospheric dynamics because the lower level of
the transition is metastable (Dupree et al. 1992, 2009; Smith et al. 2012). In
Omega Centauri, a closely matched group of first-ascent red giant stars
displays strong and weak helium absorption that correlates (Dupree et al.
2011) with increased [Al/Fe] and [Na/Fe] abundance, more than with [Fe/H].
This result gave direct observational support to the idea that products of
high-temperature hydrogen burning in a previous stellar generation had, in
fact, occurred. A quantitative measure of the helium abundance in these
objects is the goal of this Letter.

Pasquini et al. (2011) calculated profiles of the He I 1.08$\mu$m line in an
approximate way based on a stationary plane-parallel model applied to two very
cool luminous stars in NGC 2808. They showed that a change in the
chromospheric structure itself can strengthen or weaken helium absorption. In
fact, chromospheric line profiles are highly sensitive to the structure and
dynamics of the atmospheric model. In this paper, we have selected similar
stars and first constrained the atmospheric structure and dynamics using other
chromospheric lines. A model for the radiative transfer must be used that is
appropriate to the stars. Following that, the abundance of helium can be
inferred from line synthesis using the semi-empirical atmospheric model that
is anchored by other chromospheric lines.

Here we focus on two ‘identical’ red giants in Omega Centauri, LEID 54064 and
LEID 54084 (van Leeuwen et al. 2000). They are located $\sim$5.7 arc min to
the SW from the cluster center, and are separated by 1.6 arcminutes on the
sky. These giants have very similar temperatures, luminosities, and values of
[Fe/H] (Table 1). However they differ remarkably in [Na/Fe] and [Al/Fe]
abundances and the strength of the helium line (Dupree et al. 2011). The star
LEID 54084 exhibits enhanced light elements as compared to LEID 54064.

##  2 Modeling Chromospheric Lines

The PANDORA code (Avrett & Loeser 2003, 2008) is used to develop the semi-
empirical, spherical model of the chromosphere where the temperature
distribution, the turbulent velocities, and the expansion velocities are
adjusted to obtain optimum agreement between calculated profiles and
observations of chromospheric lines (H$\alpha$, H$\beta$, and Ca II-K). The
initial model consists of a static LTE photosphere corresponding to a
effective temperature of 4740K (Kurucz 2011), gravity $log~{}g$= 1.75, a
stellar radius of 20$R_{\odot}$, and [Fe/H]=$-$1.72 with the
$\alpha$-abundances enhanced by $+$0.44 dex. Chromospheric line emission is
essentially unaffected by the photospheric model. A chromospheric structure
similar to other metal-poor models (Mészáros et al. 2009) was added to begin
the iterations, and expansion started in the low chromosphere. Our
calculations assume multi-level atoms (H I:15 levels, Ca II: 5 levels, He I:
13 levels), and the iterations explicitly consider the velocity field in the
evaluation of the line source functions and as a contribution to the pressure
in the hydrostatic equilibrium equations. The total model is iterated with
full and complete non-LTE calculations in order to match the chromospheric
line profiles. The Ca II-K line profile is computed with partial frequency
redistribution; complete frequency redistribution is used for the hydrogen and
helium lines. These flux profiles are calculated with an integration over the
apparent spherical stellar disk including the extended chromosphere.

The profiles of the optical lines, H$\alpha$, H$\beta$, and Ca II-K were taken
from spectra obtained with the MIKE double echelle spectrograph (Bernstein et
al. 2003) mounted on the Magellan/CLAY telescope at Las Campanas Observatory.
These spectra were used previously to derive elemental abundances (Dupree et
al. 2011). The spectra and the calculated stellar profiles for H$\alpha$,
H$\beta$, and Ca II-K are shown in Fig. 1. The observed profiles are
effectively identical between the two giants, signaling that the activity
levels of the stars are similar. The spectra are well matched by the
calculated profiles. Note the asymmetry in the H$\alpha$ line core; the core
is formed higher in the atmosphere than the rest of the profile and is
sensitive to the outflow. However the line itself is narrow, and demands a
relatively low turbulent velocity, which increases with height in the
chromosphere. The final model (Fig. 2) has a temperature that extends to 105K
(although such high temperatures do not affect the profiles evaluated here),
and an outflow velocity that reaches 100 km s-1, which yields a mass outflow
rate of $\sim 3\times 10^{-9}$ M⊙ yr-1. This rate follows straightforwardly
from the atmospheric model (Fig. 2) and is proportional to $Nvr^{2}$ in the
chromosphere, where $r$ is the radial distance at which the wind has a
velocity $v$, and $N$ is the hydrogen density. This value exceeds by a factor
of 1.3–1.5 the rate estimated from an extension of the Mészáros et al. (2009)
fit to H$\alpha$ profiles of cooler stars shown in their Fig. 10. For more
luminous stars in the more metal-rich NGC 2808, Mauas et al. (2006) find
values of $0.7-3.8\times 10^{-9}$ M⊙ yr-1 from the H$\alpha$ line. Field
metal-poor giants, comparable in $M_{bol}$ to our targets possess a mass loss
rate spanning $1.3\times 10^{-9}-10^{-8}$ M⊙ yr-1(Dupree et al. 2009). While
mass loss rates have been measured (Mészaros et al. 2009) to vary with time by
factors of 1.5 to 6 in metal-poor red giants with luminosity
$logL/L_{\odot}\sim 3.0$, the values inferred from semi-empirical model fits
are less by an order of magnitude than the Reimers (1977), Origlia et al.
(2007), or the Schröder & Cuntz (2005) approximations.

This temperature and velocity model (Fig. 2) is used to evaluate the profile
of the He I 1.08$\mu$m line. The near-ir He I lines measured with PHOENIX on
Gemini-S were reported earlier (Dupree et al. 2011). The populations,
ionization fraction, and continuum emission, are evaluated in separate models
for each value of the helium abundance, and the profile is calculated assuming
spherical geometry in an expanding atmosphere. The contribution of the
extended chromosphere can be noted in the weak emission present on the long
wavelength side of the line. The helium absorption extends substantially
towards shorter wavelengths due to scattering in the expanding atmosphere and
is enhanced by the metastable nature of the lower level of the transition. The
helium lines are essentially P Cygni profiles since the red giants have
extended atmospheres. The population of the lower level of the 1.08$\mu$m
transition peaks at T=18,000K, but lies within a factor of two of its maximum
value between 14,000 and 25,000K; the outflow velocity doubles over this
temperature span.

Various values of the helium abundance, from Y=0.15 to Y=0.50 [log(nHe/nH)
ranging from 10.65 to 11.4], were assumed and 9 models calculated. The
abundance selected minimizes the residuals between the observed and calculated
profiles. The star LEID 54084 clearly exhibits a broad helium line which could
extend to shorter wavelengths beyond the Si I absorption at 1.027$\mu$m but is
compromised by the presence of the water vapor blend with Si I. A value of
Y=0.39 to Y=0.44 well represents the depth of the observed profile
representing the minimum range in the residuals. Helium is not clearly
detected in LEID 54064. The calculated profiles for Y$\leq$0.22 give a minimum
in the residuals, and we adopt this value as an upper limit to Y. Inspection
of the helium profiles shows that a value of Y=0.25 overpredicts the strength
of the line in LEID 54064, and the residuals of the fit are larger than for
Y=0.22. These simulations suggest that the helium abundance difference is
$\Delta Y\geq 0.17$ between the two stars.

##  3 Discussion and Conclusions

This spectroscopic value of helium from LEID 54084, namely Y=0.39–0.44 can be
compared to values obtained from models of stellar structure and evolution. In
Omega Cen, Norris (2004) estimated the presence of helium from isochrones
matching the lower main sequence with values of Y ranging from 0.23 to 0.38.
Piotto et al. (2005) noted the blue main sequence could only be matched with
stellar models with helium abundance ranging from 0.35 $<$ Y $<$ 0.45 and
concluded that Y=0.38 best fit the ridgelines in the color-magnitude diagram
of Omega Cen. HST photometry of an outer field in the cluster (King et al.
2012), reveals a helium abundance for the blue main sequence of
Y=0.39$\pm$0.02. Recent Yonsei-Yale isochrones for several subpopulations in
Omega Cen (Joo & Lee 2013) suggest a range in Y from 0.38 to 0.41. Thus the
spectroscopic value of helium for LEID 54084, a star with enhanced light
element abundances is in harmony with the abundance inferred from stellar
structure models. In a more metal rich cluster, NGC 2808, the approximate
model of Pasquini et al (2011) suggested one star may have a similar value of
Y=0.39 to 0.5.

The Y value for LEID 54064 where the helium line is weak (or not detected) has
an upper limit (Y $\leq$ 0.22) that is slightly less than the cosmic value
(Y=0.24). These abundances suggest the helium enhancement, $\Delta$Y, is
$\geq$0.17. King et al. (2012) concluded from plausible fits to the color
magnitude diagram of Omega Cen that $\Delta$Y$\sim$0.15 where a value for the
primeval abundance of helium (Y=0.24) was chosen for the red main sequence.
Piotto et al. (2005) required $\Delta$Y=0.14 to explain the differences in
metal abundances found for the blue and red main sequences. It is interesting
to note that the Sun requires a helium abundance of Y=0.27$-$0.28 to match the
solar luminosity, but due to diffusion and settling, the helium abundance in
the envelope is less, Y=0.24$-$0.25 (Christensen-Dalsgaard 2002; Guzik & Cox
1993), and Y=0.16 (corresponding to nHe/nH=0.05) in the steady-state solar
wind (Kasper et al. 2007). It may be that spectroscopy will yield different
values for the helium abundance from those inferred from stellar isochrone
models, although currently we do not know if the characteristics of the solar
abundance pattern occur in these metal-poor giant stars.

The optical and near-infrared spectra used here were acquired about 3 months
apart, and a variation in the line profiles might occur. However, these giants
have log $L/L_{\odot}~{}\sim$ 2.2, and $M_{V}\sim-$0.45 and lie on the red
giant branch below the stars that exhibit H$\alpha$ wing emission. It is this
emission which can vary in strength in first-ascent red giants (Mészáros et
al. 2008; Cacciari et al. 2004). The remarkable similarity of H$\alpha$,
H$\beta$, and Ca II-K profiles between the two giants suggests that activity
does not cause significant changes.

Another consideration might be the presence of X-rays or EUV emission from a
high temperature plasma. Because neutral helium can be photoionized and then
recombine preferentially into the lower level of the 1.08$\mu$m line, this
process would enhance the strength of the observed helium line. Red giants
need substantial magnetic confinement of material to produce hot plasma;
magnetic signatures in the spectra of similar stars have not been detected,
and the coronae appear absent (Rosner et al. 1995). The slightly metal-poor K
giant, $\alpha$ Boo, has a ‘tentative detection’ (Ayres et al. 2003) of X-rays
but, if indeed present, they are a factor of 104 weaker in $L_{X}/L_{bol}$
than the average solar value and would seem to have little effect on the
profile.111CHANDRA images of Omega Cen (Cool et al. 2013) do not reach faint
sources ($L_{X}\lesssim 10^{29}\ erg\ s^{-1}$). The identified optical
counterparts of the X-ray sources are binaries, and not the single red giants
that are targeted here. In $\alpha$ Boo, the equivalent width of the
1.083$\mu$m line varies in absorption strength which could be caused by wind
variation as well as chromospheric excitation conditions (O’Brien & Lambert
1986). Single metal-poor red giants in the field also display a very weak
1.08$\mu$m absorption line, and though these stars are optically brighter,
they have not been detected in X-rays. Population I giants, which generally
exhibit X-rays, have stronger helium absorption as compared to their metal-
poor field counterparts (cf. Dupree et al. 2009). This suggests the line is
not influenced by X-rays in the metal-poor stars. The Ca II-K lines are very
similar in the two giants (Fig. 1) indicating that these stars have similar
chromospheres such that X-rays would not be present in only one star causing
the strengthening of the helium absorption. Thus it does not appear likely
that X-rays contribute to the line formation for the targets considered here.
Several epochs of measurement would clearly be useful to determine if
variation occurs in the helium lines.

Pasquini et al. (2011) carried out a similar calculation for two stars in the
globular cluster NGC 2808. The 2 luminous stars ($logL/L_{\odot}\ \sim 3.2$)
selected by Pasquini et al. (2011) have different levels of activity as
indicated by the Ca K line which underscores the ubiquitous variability of
such luminous giants. These differences demand different semi-empirical models
for the two stars yet only one model was used; in addition the observations of
the optical and infrared spectra were separated by some weeks which brings
uncertainty when modeling such luminous active objects. Computation of the
line profiles in Pasquini et al. invokes models that do not adequately
represent the stars nor the conditions in their atmospheres. The use of a
plane-parallel approximation is questionable when modeling a star of radius
$\sim$84 R⊙. The computation assumed a static atmosphere, and the authors
simply shifted the calculated line in wavelength to match observations.
However, the spectra show that the chromosphere, as measured by H$\alpha$, Ca
K, and the helium line, exhibits signatures of outflow. We have taken our
model and calculated the helium profile under the same assumptions adopted by
Pasquini et al. (2011) (plane parallel and static) for comparison to a model
with the appropriate assumptions for these stars, namely spherical geometry
and expanding. The results of this calculation show substantial differences.
Not only does the spherical model exhibit emission, but the absorption is
larger than the static model due to the expanding atmosphere. A larger star,
with extended chromosphere and/or wind, might be expected to exhibit more
substantial changes. For the same value of the helium abundance, the
equivalent width of the absorption in the expanding spherical model is larger
by 5 to 19% than the static plane-parallel model depending on the value of Y.
(Here we assumed Y=0.28 and Y=0.44.) Thus, interpreting the observed profile
formed in an expanding large giant star, by ’matching it’ to a static, plane
parallel profile, as did Pasquini et al. (2011), will lead to an overestimate
of the abundance of helium. Pasquini et al. (2011) do not compare the computed
profiles of Ca K and H$\alpha$, to the stellar spectra so the adequacy of the
models is unknown. Consideration of all of these facts indicates that the
determination of the helium abundance in Pasquini et al. (2011) must be
approached with caution.

The targets selected in this paper are of much lower luminosity where
variability is absent or greatly minimized. Moreover, the 2 stars are
effectively identical in temperature, luminosity, iron abundance, activity,
and in chromospheric features - with the exception of helium and enhanced Al
and Mg. The treatment of the radiative transfer is state of the art with a
spherical atmosphere, assuming an outflow, where the outflow is incorporated
into the source function for the lines.

The abundance of helium and its variation between these two giant stars in
Omega Cen gives quantitative observational confirmation of a helium
enhancement to accompany the enhanced light metals. The near-ir line of He I
can provide a probe of the helium abundance in cool stars when additional
chromospheric profiles are available to constrain the atmospheric structure
and dynamics and appropriate radiative transfer calculations are employed.

We are grateful to Bob Kurucz who calculated specific photospheric models to
initiate the calculations. Facilities: Gemini:South (PHOENIX), Magellan:Clay
(MIKE)

## References

  * (1)
  * (2) Alonso, A., Arribas, S., & Martinez-Roger, C. 1999, A&AS, 140, 261
  * (3)
  * (4) Anderson, J. 1997, PhD thesis, Univ. California, Berkeley
  * (5)
  * (6) Avrett, E. H., & Loeser, R. 2003, in IAU Symp. 210, Modeling of Stellar Atmospheres, ed. W. Weise & N. Piskunov (Dordrecht: Kluwer), A-21
  * (7)
  * (8) Avrett, E. H., & Loeser, R. 2008, ApJS, 175, 229
  * (9)
  * (10) Ayres, T., Brown, A., & Harper, G. M. 2003, ApJ, 598, 610
  * (11)
  * (12) Bedin, L. R., Piotto, G., Anderson, J., et al. 2004, ApJ, 605, L125
  * (13)
  * (14) Bellini, A., Bedin, L. R., Piotto, G., et al. 2010, AJ, 140, 631
  * (15)
  * (16) Bernstein, R. A., Shectman, S. A., Gunnels, S. M., Mochnacki, S., & Athey, A. E. 2003, Proc. SPIE, 4841, 1694
  * (17)
  * (18) Brodie, J. P., & Strader, J. 2006, ARA&A, 44, 193
  * (19)
  * (20) Cacciari, C., Bragaglia, A., Rossetti, E. et al. 2004, A&A, 413, 343
  * (21)
  * (22) Charbonnel, C., Krause, M., Decressin, T., Prantzos, N., & Meynet, G. 2013, arXiv 1301.5016
  * (23)
  * (24) Christensen-Dalsgaard, J. 2002, Rev. Mod. Phys., 74, 1073
  * (25)
  * (26) Cool, A. M., Haggard, D., Arias, T., et al. 2013, ApJ, 763, 126
  * (27)
  * (28) DaCosta, G. S., Norris, J., & Villumsen, J. V. 1986, ApJ, 308, 743
  * (29)
  * (30) deMink, S. E., Pols, O. R., Langer, N., & Izzard, R. G. 2009, A&A, 507, L1
  * (31)
  * (32) D’Ercole, A., D’antona, F., Ventura, P., Vesperini, E., & McMillan. S. L. W. 2010, MNRAS, 407, 854
  * (33)
  * (34) Dupree, A. K., Sasselov, D. D. & Lester, J. B. 1992, ApJ, 387, L85
  * (35)
  * (36) Dupree, A. K., Smith, G. H., & Strader, J. 2009, AJ, 138, 1485
  * (37)
  * (38) Dupree, A. K., Strader, J., & Smith, G. H. 2011, ApJ, 728, 155
  * (39)
  * (40) Geyer, E. H. 1967, ZAp, 66, 16
  * (41)
  * (42) Gratton, R. G., Carretta, E., & Bragaglia, A. 2012, A&A Rev., 20, 50
  * (43)
  * (44) Guzik, J. A., & Cox, A. N. 1993, ApJ, 411, 394
  * (45)
  * (46) Johnson, C. J., & Pilachowski, C. A. 2010, ApJ, 722, 1373
  * (47)
  * (48) Joo, S. J., & Lee, Y. W. 2013, ApJ, 762, 36
  * (49)
  * (50) Kasper, J.C., Stevens, M. L., Lazarus, A. J., Steinberg, J. T., & Ogilvie, K. W. 2007, ApJ, 660, 901
  * (51)
  * (52) King, I. R., Bedin, L. R., Cassisi, S., et al. 2012, AJ, 144, 5
  * (53)
  * (54) Kurucz, R. L. 2011, Can. J. Phys., 89, 417
  * (55)
  * (56) Martell, S. L. 2011, AN, 332, 467
  * (57)
  * (58) Mauas, P. J. D., Cacciari, C., & Pasquini, L. 2006, A&A, 454, 609
  * (59)
  * (60) Mészáros, Sz., Dupree, A. K., & Szentgyorgyi, A. 2008, AJ, 135, 1117
  * (61)
  * (62) Mészáros, Sz., Avrett, E. H., & Dupree, A. K. 2009, AJ, 138, 615
  * (63)
  * (64) Moehler, S., Dreizler, S., Lanz, T., et al. 2011, A&A, 526, A136
  * (65)
  * (66) Moni Bidin, C., Villanova, S., Piotto, G., et al. 2012, A&A, 547, A109,
  * (67)
  * (68) Norris, J. E. 2004, ApJ, 612, L25
  * (69)
  * (70) O’Brien, G. T., & Lambert, D. L. 1986, ApJS, 62, 899
  * (71)
  * (72) Origlia, L, Rood, R. T., Fabbri, S. et al. 2007, ApJ, 667, L85
  * (73)
  * (74) Pasquini, L., Mauas, P., Käufl, J. U., & Cacciari, C. 2011, A&A, 531, A35
  * (75)
  * (76) Piotto, G., Villanova, S., Bedin, L. R., et al. 2005, ApJ, 621, 777
  * (77)
  * (78) Reimers, D. 1975, Mem. Soc. R. Sci Liége, 8, 369
  * (79)
  * (80) Renzini, A. 2013, Mem. S. A. It., in press (arXiV:1302.0329)
  * (81)
  * (82) Rosner, R., Musielak, Z. E., Cattaneo, F., Moore, R. L., & Suess, S. T. 1995, ApJ, 442, L25
  * (83)
  * (84) Schröder, K.-P., & Cuntz, M. 2005, ApJ, 630, L73
  * (85)
  * (86) Skrutskie, M. F., Cutrie, R. M., Stiening, R. et al. 2006, AJ, 131, 1163
  * (87)
  * (88) Smith, G. H., Dupree, A. K., & Strader, J. 2012, PASP, 124, 1252
  * (89)
  * (90) Valcarce, A. A. R., Catelan, M. & Sweigart, A. V. 2012, A&A, 547, A5
  * (91)
  * (92) van Leeuwen, F., Le Poole, R. S., Reijns, R. A., Freeman, K. C., & de Zeeuw, P. T. 2000, A&A, 360, 472
  * (93)
  * (94) Woolley, r. v. d. R. 1966, R. Obs. Ann., 2,1
  * (95)

Figure 1: H$\alpha$ , H$\beta$, and Ca II K-line region shown for the two
giant stars with the model fit overlaid (indicated by a solid line which is
colored red in the electronic edition). The spectra of the two stars are
virtually identical. The Ca II spectrum displays an interstellar absorption
feature blended with the Fe I line at $-$3Å from the line core in the LEID
54064 spectrum, but distinct from the Fe I feature in the LEID 54084 spectrum.
Weak emission may be present on the long-wavelength wing of the H$\alpha$ line
and on both wings of the H$\beta$ line. The success of the model can be seen
from the agreement between observed and calculated chromospheric profiles. A
color version of this figure is in the electronic edition. Figure 2: Final
model. The top panel displays the total hydrogen density (left axis) and the
temperature (right axis). The lower panel shows the turbulent and outflow
velocities needed to match the observed profiles.

Figure 3: Observed Helium lines binned to a resolution element in the two
matched red giants with the 10830Å profile as calculated for several values of
the helium abundance (shown by a solid smoothly varying line that is colored
red in the electronic edition). Values of Y are given for the calculated
curves arranged top to bottom and the corresponding log (nHe/nH) is shown in
parentheses where log $n_{H}$=12.00.  Table 1: Characteristics of Target Stars

Quantity | LEID 54064 | LEID 54084 | Refs.
---|---|---|---
V | 13.27 | 13.21 | 1
B$-$V | 1.048 | 1.044 | 1
Ks | 10.62 | 10.56 | 2
Teff [K] | 4741 | 4745 | 3
log g [cm s-2] | 1.76 | 1.74 | 3
MV |  $-$0.43 |  $-$0.49 | 4
log $L/L_{\odot}$ | 2.21 | 2.23 | 5
$[Fe/H]$ |  $-$1.86 |  $-$1.79 | 3
$[Na/Fe]$ |  $-$0.14 | 0.37 | 3
$[Al/Fe]$ |  $\leq$0.36 | 1.12 | 3
EW (He I) [mÅ] |  $\leq$9.2 | 89.5 | 3

References. — (1) van Leeuwen et al. 2000 ;(2) 2MASS All Sky Survey; Skrutskie
et al. 2006; (3) Dupree et al. 2011; (4) Distance modulus from Johnson &
Pilchowski 2010; (5) Bolometric correction from Alonso et al. 1999.