Helium in the galaxy

Yistas in Astronomy, Voi.24, pp.355-373. ©Pergamon Press Ltd, 1981. Printed in Great Britain. Editors: A. Beer, K. Pounds and P. Beer.

HELIUM

0083-6656/81/0301-0355505.00/0

IN THE G A L A X Y C. T h u m

Max-Planck-Institut for Astrophysik, D-8046 Garching, FRG

I.

Introduction The astronomy of radio recombination lines began with Kardashev's (1959) prediction

that l i n e s a r i s i n g from t r a n s i t i o n s between high energy levels of recombining ions in ionized i n t e r s t e l l a r clouds ( i . e . HII regions) may be observable. The subsequent detection of such radio recombination l i n e s , at wavelengths of 6 and 18 cm (H~glund and Mezger 1965; L i l l e y et a l . 1966a) from hydrogen in the b r i g h t galactic HII regions Orion A and MI7 marked the beginning of a f r u i t f u l

branch of g a l a c t i c astronomy which is not yet f u l l y exploited.

Shortly t h e r e a f t e r , astronomers at Harvard Observatory published the detection of a recombination l i n e of helium, the second most abundant element in i n t e r s t e l l a r space ( L i l l e y et a l . 1966b). The s c i e n t i f i c potential of t h i s discovery was realized immediately. F i r s t l y , through an accurate determination of the helium-to-hydrogen l i n e i n t e n s i t y r a t i o , i t is possible to measure the abundance of helium in ionized regions of the i n t e r s t e l l a r medium. Secondly, by comparing the helium and hydrogen l i n e widths, thermal and t u r b u l e n t v e l o c i t i e s can in p r i n c i p l e be separated in a nebula, and thus a s t r a i g h t forward determination of i t s electron temperature appears possible. Although the l a t t e r did not prove successful, the first

aspect has motivated much of the subsequent research. Helium abundances in Galactic emission nebulae can also be determined from optical spec-

troscopy, although the l i n e most often used (~5876) is r e l a t i v e l y weak. I t is indicated in the He energy level diagram (Fig. i ) together with a few even weaker t r a n s i t i o n s occasionally observed also. The strong resonance l i n e at 584 R is not observable because of very strong absorption by i n t e r s t e l l a r gas. In deriving the abundance of (ionized) helium from optical l i n e s the 5876 ~ l i n e is usually compared with a nearby t r a n s i t i o n , often HB (n = 4 ÷ 2) of hydrogen whose energy levels are also displayed in Fig. i . Given the substantial differences in the level diagrams of both elements i t is clear that t h i s comparison is not s t r a i g h t f o r ward, e s p e c i a l l y i f c o l l i s i o n a l processes are involved. In f a c t , a model of the l i n e - e m i t t i n g region and a f a i r l y detailed knowledge of the helium e x c i t a t i o n are required. This problem becomes less severe at higher quantum numbers n, where the energy levels of He and H are more s i m i l a r . In the radio regime (~>I mm, n >38), the level diagrams are i n d i s t i n g u i s h a b l e and, as a consequence, the i n t e n s i t y r a t i o of corresponding He and H t r a n s i t i o n s should immediately y i e l d the ionized helium abundance. I t s determination from radio data is thus considered to be more r e l i a b l e than from optical data. Another advantage is that radio observations are not hindered by i n t e r s t e l l a r dust. This is of crucial importance, since the vast majority of the emission nebulae l i e in the g a l a c t i c plane where e x t i n c t i o n is highest, thus severely l i m i t i n g the o p t i c a l l y accessible portion of 355

H

356

C. Thum

s,p,d....

s

p

the g a l a c t i c disk. Also, dust w i t h i n the HII

Hel

regions themselves renders u n c e r t a i n the o p t i -

f

d

n=10 9

eV 0.2 --

cal data on i n d i v i d u a l sources. Therefore, •m

8

eV 03

m.m.m.i

radio observations of He and H recombination l i n e s o f f e r a unique o p p o r t u n i t y to determine

7

the abundance of helium throughout the Galaxy.

6

....

5

....

4`

. . . .

0.5

The compactness of the b r i g h t HII regions

0.5 ---"

- -

and the weakness of the l i n e s makes the study

Z

of recombination l i n e s the domain of the big

I

1.

/ /'

C~

.

.

.

radio telescopes. Most of the work has been done w i t h the 43-m telescope at Green Bank,

6

.

2.

/ // ,

'\

.

USA, the 64-m telescope at Parkes, A u s t r a l i a

k

and, since 1973, the lO0-m telescope at Effels.

'

berg, Germany. Fig. 2 shows a t y p i c a l recombi-

_7/. _

5.

nation l i n e spectrum obtained w i t h the lO0-m 5.

telescope at a wavelength of 6 cm. The main f e a t u r e is the hydrogen recombination l i n e

584k

10. 13.6

HlO9a accompanied by the corresponding helium

lO.

and carbon t r a n s i t i o n s ,

P

as well as the H1376

l i n e . Here the na l i ~ e s r e f e r to t r a n s i t i o n s 20. 24,.6

n=1

20.

between n + i and n, nB l i n e s to t r a n s i t i o n s n + 2, etc.

l,

s

Is,p,d ....

p

d

f

The frequencies o f these t r a n s i t i o n s are

HeI

H

c a l c u l a t e d from the Rydberg formula

Fi~. I. Level diagram of hydrogen (left) and neutral helium (right). Singlet states of helium are dashed, whereas the triplets are shown as solid lines. The strongest optical line (5876 ~) from He in HII regions is indicated together with some weaker optical transitions (dashed-dotted). The strong resonance transition (584 ~) is not observable (Schmid-Burgk 1980). i

I

1.0-

'

'

'

I

'

'

'

I

_

cR

v

'

'

I

'

~

Z2

/

---me 1+ emA_m

'

'

I

1

1

t n2

'

'

m2

'

HlOga

l..--

0.5

Z

HelOga C109(x /

0.0 I

I

-200

I

I

I

I

-lO0

i

L

/ J

I

I

)

(I.i)

'

'

W ~8

/

J

i

i

I

i

0 100 RADIALVELOCITY VLSR[KM S-1]

l

L

I

200

i

i

l

I

J

l

300

Fi@. 2. Recombination line spectrum at 5 GHz towards W48 obtained with the lO0-m telescope (Churchwell et al., unpublished). The velocity scale refers to the HIOga line and is relative to the local standard of rest.

Helium in the Galaxy

357

where c is the v e l o c i t y of l i g h t , R~ = 1.0973732 108 cm-1 is the Rydberg constant for i n f i n i t e mass, me and mA are the masses of the electron and the neutral atom, Z the e f f e c t i v e nuclear charge, m and n are the upper and lower quantum levels involved. Due to t h e i r d i f f e r e n t atomic masses He lines occur at frequencies about ~0.04 percent higher than the H-transition at the same n which corresponds to a displacement of -122.2 km s - I on the hydrogen v e l o c i t y scale. In normal HII regions t h i s is enough to separate the two lines c l e a r l y , since t h e i r h a l f power widths are of the order of 30 km s- I . Spectra l i k e the one in Fig. 2 have been used in the past decade to derive a picture of the d i s t r i b u t i o n of helium in the Galaxy. I t is the aim of t h i s paper to present t h i s picture together with the main milestones along the way. This subject has been reviewed before in the more general context of radio recombination lines by Mezger (1976) and Brown et a l . (1978). I t is included in Panagia's (1979) paper on galactic chemical abundances, and is s p e c i f i c a l l y addressed by Mezger (1980). 2.

Determination of the lonized Helium Abundance In radio recombination l i n e work, the basic quantity observed with a telescope whose

antenna pattern, main beam solid angle and beam e f f i c i e n c y are f ( ~ ) , ~m and n B, respectively, is the frequency dependent antenna temperature TA, L and i t s integral ~ ( i ) : n B /4_

r_

c

Here TB, L is the brightness temperature of the l i n e which is position-dependent. Helium abundances are calculated from the r a t i o of the integrated He and H l i n e p r o f i l e s as given by eq. (2.1). ~ ( i ) depends on telescope and receiver characteristics which change with frequency. This PrOblem is solved by observing Hen~ and Hn~ lines whose frequency difference (0.04 percent) is n e g l i g i b l e in t h i s context. At higher frequencies, where the received signals c r i t i c a l l y depend on elevation and atmospheric conditions, both lines should be observed simultaneously, so that time- and elevation-dependent effects of the atmosphere, telescope gain and receiver c a l i b r a t i o n cancel. I f the recombination l i n e considered is o p t i c a l l y t h i n , as is nearly always the case, the l i n e brightness temperature, under LTE conditions, is the product of the optical depth TL and the electron temperature Te along the l i n e of sight s. Since the l i n e absorption c o e f f i c i e n t KL is proportional to nen(i)T e-2.5 (e.g. Mezger 1974) we have

TB'L

= ~Te dTL ~ fnen(i)Te-l"5s ds

(2.2)

where ne and n ( i ) are the electron and ion densities, respectively. Combining eqs. (2.1) and (2.2) gives -Jnen(i)Te - I ' 5 dV

(2.3)

V(i)

where now the integration is performed over the volume occupied by the recombining ion i in the telescope beam. For a constant electron temperature eq. (2.3) y i e l d s the abundance (by number) of ionized helium y+ as the r a t i o of the integrated He and H l i n e p r o f i l e s

C. Thum

358

e(He+) e(H +)

:

n(He+) n(H +)

= y+

(2.4)

Observations indicate that Te varies in an HII region, and that these v a r i a t i o n s may be correlated with density v a r i a t i o n s . Compared with density v a r i a t i o n s , which are known to range t y p i c a l l y over more than an order of magnitude, the v a r i a t i o n of Te in a given HII region appears to be quite small. For example, v a r i a t i o n s of less than 30 percent are observed in the Orion Nebula (Pankonin et a l . 1980). Constancy of Te is t h e o r e t i c a l l y understandable because cooling by forbidden l i n e s tends to s t a b i l i z e Te at values ~i04 K. Constant Te is therefore a rather safe assumption in t h i s context. Another assumption is more c r i t i c a l ,

however. The simple p r o p o r t i o n a l i t y in eq. (2.2)

presumes that the energy levels at the recombining atoms are populated with a Boltzmann d i s t r i b u t i o n , i . e . they are in local thermodynamic e q u i l i b r i u m (LTE). However, deviations of the level populations from t h e i r LTE values, described by m u l t i p l i c a t i o n factors bn, can be appreciable because real HII regions are by no means perfect black bodies. In a d d i t i o n , the r e l a t i v e l y strong f r e e - f r e e continuum r a d i a t i o n f i e l d present in an HII region ( t y p i c a l l i n e to-continuum r a t i o s are 0.05 at 5 GHz) can increase the l i n e strength by stimulated emission. Allowing f o r these e f f e c t s , the l i n e brightness temperature can be expressed (Dupree and Goldberg 1970) as %c * bn (1 + Y (I - ~n) ) TB,L = TB,L

which holds for continuum optical depths mc <<1 •

(2 5)

T * is the LTE value obtained from eq. B,L

(2.2). The function Bn = I - kTe d(In bn) hv dn

(2.6)

describing the stimulated emission is tabulated together with the bn values by Brocklehurst (1970) f o r a range of electron densities and temperatures. For a general s o l u t i o n of the t r a n s f e r of recombination l i n e r a d i a t i o n including non-LTE effects and pressure broadening, the reader is referred to the classic paper by Brocklehurst and Seaton (1972). I t is e a s i l y seen from eq. (2.6) that Bn can a t t a i n large negative values even i f the bn factors are close to u n i t y and vary slowly with n. TB, L can thus be increased considerably over i t s LTE value. At higher frequencies (~15 GHz), however, where Tc <10 -3 in the m a j o r i t y of HII regions, non-LTE effects are small. The c r i t i c a l

frequency above which non-LTE effects

can be neglected depends on the emission measure of the source and i t s density s t r u c t u r e , and is not e a s i l y determined (see Walmsley 1980). To assess the importance of deviations from the LTE approximation (eq. 2.4), i t is necessary to carry out a numerical i n t e g r a t i o n of eq. (2.1) substituted with (2.5) and (2.6) for an adopted density model. This has been done f o r the Orion Nebula (Batchelor 1974) f o r the density model suggested by Brocklehurst and Seaton (1972) For frequencies of 5 GHz or higher (n <109) the numerical s o l u t i o n does not deviate from the LTE r e s u l t , but with decreasing frequency opacity effects dominate in the central (densest) part of the nebula and non-LTE effects dominate in the outer layers. In combination, the two effects y i e l d the important r e s u l t , with very few exceptions,

Helium in the Galaxy

~(He+)/~(H +) ~y+

359

(2.7)

Although t h i s r e s u l t has e x p l i c i t e l y been derived for the Orion Nebula density d i s t r i b u t i o n , i t is expected to be generally v a l i d . Observations show that giant HII regions are composed of a number of compact components imbedded in a low density envelope. Hence, for density d i s t r i b u t i o n s resembling that of the Orion Nebula, eq. (2.7) can be assumed to hold for n ~i09. At higher n the envelope contributes strongly to the l i n e i n t e n s i t y , and the r e s u l t w i l l depend on the detailed density d i s t r i b u t i o n . At any rate, two statements can be made (Churchwell, Mezger and Huchtmeier 1974; hereafter Paper I ) :

( i ) I f the volumes V(i) in eq. (2.3) are the same for He and H, the He/H l i n e r a t i o

equals the helium abundance y regardless of non-LTE e f f e c t s , i . e . E(He+)/E(H +) = y (ii)

I f V(He+)~V(H+) which is always the case for e x c i t a t i o n by OB stars (section 3), we have E(He+)/E(H +) ~y

(2.8)

This means that non-LTE e f f e c t s , Stark broadening and opacity only a f f e c t the l i n e r a t i o i f He+ and H+ volumes do not coincide. Unfortunately, t h i s is expected to be the case as we shall see in the next sections. Even then, however, eq. (2.7) and (2.8) allow firm conclusions to be drawn from observations. Among the effects which cause deviations from the simple LTE s o l u t i o n (eq. 2.4) s t i m u l a t ed emission is thought to be the most serious one. Opacity effects which depend on frequency as v - 2 " I can be avoided i f observations are made at frequencies about 3 times higher than the turn-over frequency in continuum spectrum. Pressure broadening has an even stronger frequency dependance (&~-v-2"5). This e f f e c t is probably n e g l i g i b l e at quantum numbers n ~i09 except perhaps f o r the densest HII regions ( P i t a u l t and Cesarsky 1980) and may lead to an underestimate of the observed l i n e strengths, since some f r a c t i o n of the l i n e i n t e n s i t y is s h i f t e d into broad wings which are d i f f i c u l t

to observe. Because the He l i n e probably o r i g i n a t e s in

the densest parts of the HII region, i t w i l l be more affected than the H l i n e , and t h i s + process w i l l r e s u l t in an underestimate of y . Stimulated emission depends on the continuum optical depth (eq. 2.5) and i t s importance is therefore reduced at higher frequencies. A detailed i n v e s t i g a t i o n of non-LTE effects is given by Brown et a l . 1978. Starting from a model HII region with a steep radial density gradient the authors predict considerable deviations of observable q u a n t i t i e s from t h e i r LTE values at nearly a l l radio frequencies. While these strong non-LTE effects do not appear to be observed in real HII regions (Shaver and Wilson 1979; Walmsley 1980; Shafer 1980) these model c a l c u l a t i o n s reinforce the conclusion that the measured He/H l i n e r a t i o is a r e l i a b l e lower l i m i t to the He abundance. In summary, i t appears that the adverse effects of opacity, pressure broadening and stimulated emission can be avoided i f observations are done at a high enough frequency. For a typical radio HII region t h i s usually means quantum numbers n ~I09, i . e . ~ m 5 GHz, at which frequency most of the work has been done in the past.

360

C. Thum

3.

The Helium l o n i z a t i o n Structure A determination of the t o t a l helium abundance y in an HII region requires knowledge of

the abundance in a l l i o n i z a t i o n states Y = yO + y+ + y++

(3.1)

Only y+ and y++ can be determined from radio recombination l i n e s as described in section 2. Neutral helium is inaccessible to radio and o p t i c a l observations because i t is hardly e x c i t ed (Fig. 1). He+

has an i o n i z a t i o n potential of 54 eV. Normal main sequence stars do not emit a

s i g n i f i c a n t number of photons with energies higher than 54 eV. He++ is therefore not expected in HII regions, and, in f a c t , s e n s i t i v e measurements give upper l i m i t s of y++ <0.003 f o r some b r i g h t e r HII regions (Palmer et a l . 1969; and paper I ) . The e x c i t i n g stars of planetary nebulae, however, can have considerably higher temperatures than main sequence 0 stars, and a substantial f r a c t i o n of He in planetary nebulae is expected to be doubly ionized. He++ produces recombination l i n e s in a s i m i l a r way as H+ and He+ but at a higher f r e quency because Z(He+) = 2 (eq. i . i ) .

To minimize observational problems, He+n'~, Hen and

Hn l i n e s nearby in frequency are observed, so that n' =22/3n. At 23 GHz, for example, one observed He+105a together with He66~ and H66~ . Fig. 3 shows these lines towards NGC7027, the brightest planetary nebula at radio wavelengths. I t is seen that the He and He+ lines are

~66a

comparatively strong. A gaussian least squares analysis gives y+ = 0.059 and y++ = 0.049. In deriving the He++ abundance the relation '

0.3

'

I

....

I''//'

'~

o.

1'''

' I ,'

, ' I ,,,

0.2

++ I Y = 4 r.....L¢(He++)Ic(H+)j .

is used (paper I ) . The abundances obtained are in f u l l agreement with optical determina-

0.I

tions (Pankonin et al. 1981). Helium lines from other planetaries could not have been

A

0.0

detected yet because of lack of s e n s i t i v i t y al

l , l l a

I i I / / I

0

Ii

Ill

50

d L U LII

0

I

50

II

(Walmsley et al. 1981). Since these objects

I00

are quite small and generally outside the ''

I ....

I''

'//'

''I

''''I,,

,'I'

0.2

Galactic plane, they are better observed at optical wavelengths.

0.1

ionized or neutral. I t is customary to measure

In HII regions helium is only singly

He+105a

the fraction of neutral helium from the quant i t y R defined as the ratio of the He+ and H+

0.0

.

L

I

I , , , ,

0

I,

i

n / / , ,

,

J

50 0 VELOCITY (KM S-I]

.

,

.

volumes weighted by the proton density squared

.

,

.

i

I

50

,

,

,

i

I

,

a

100

Fi~. 3. Recombination line spectrum at 22.4 GHz towards NGC7027 obtained with the lO0-m telescope (Pankonin, Thum and Mezger, in prep.). Gaussian curves (dashed) are fitted to the observed lines.

R = f n2 (H+)dV/fn 2 (H+)dV V(He +)

(3.2)

V(H +)

Although this quantity is hard to determine observationally, i t is c r i t i c a l for determining the total He abundance from y+. Since in

Helium in the Galaxy

361

actual observations the telescope beamwidth 0A is often smaller than the size 0s of the source, integration in eq. (3.2) is restricted to the source volumes V'(He+) and V'(H+) within the telescope beam, thus yielding the quantity R° relevant for observations.

= f n2 (H+)dV/fn 2 (H+)dV Ro v,(He +) v,(H +)

(3.3)

Inserting eq. (3.3) into (3.1) and setting I

1.0

I

I

I

I

I

I

I

y

I

++

= O, y <<1 and Ro ~1, we d e r i v e (paper I)

Ro y

0.8

+

-

(1 + y)yRo l+YRo

0.6

~y R o

(3.4)

The quantity R° has been calculated for the Orion Nebula density distribution (Batchelor

0.4

~

0

.

3

2

3

-

-

-

-

1974). Fig. 4 shows Ro as a function of the

-

resolution parameter ~ = OA/@s for various values of R. Only for complete coincidence of

0.2 0.0

L

He+ and H+ volumes (R = R° = 1), is Ro inde-

O.01L I

,

0.2

I

I

0.4

I

I

0.6

i

i

0.8

i

1.0

pendent of resolution. I f He+ f i l l s

only the

(inner) 50 percent of the (weighted) H+ volume, R rises from R, i t s value at low resolution, o

F ~ . 4. Quantity R 0 (eq. 3.3 and text) as a function of telescope resolution ~ = P A for the Orion Nebula (Batchelor 1974). Os 0A is the telescope HPBW, 0 s is the HPBW of the source brightness distribution. The curve parameter R is defined in eq. (3.2) and text.

to a maximum at f a i r l y high resolution (0A <0.20 Os). Similar results w i l l hold i f the density d i s t r i b u t i o n is homogeneous or c e n t r a l l y peaked.

Eq. (3.4) and Fig. 4 demonstrate that measurements w i l l give higher y+ values with increasing resolution, and thus a better approximation to the total He abundance. I t is this geometric e f f e c t along with the considerations presented in section 2 (opacity, pressure broadening, stimulated emission) which forces observers to higher frequencies where spatial resolution is higher. For typical radio HII regions (@s ~2 arc min) resolutions of 0A
362

C. Thum

I

I

/

Fi~. 5. Parameter R o (eq. 3.2) measuring the coincidence of He + and H + volumes as a function of the ratio X of the helium and hydrogen ionizing stellar photons (Mathis (:971).

I

/ / /

1.0

/./

is approached assymptotically for y >0.15. This corresponds to a spectral type of 09 i f

T0.5

Auer and Mihalas (1972) model atmospheres are adopted (but see section 5). Most giant HII regions require e x c i t a t i o n by more than

o oo.

I

I

I

0.05

0.1 G----,,.

0.15

one luminous 0 star (paper I ) , and one might

0.2

expect that the population of these young stars is well represented by the I n i t i a l Mass Function (Salpeter 1955). In this case,

06 stars would contribute most to the ionization of galactic HII regions, and about 90 percent of the ionizing radiation would come from stars e a r l i e r than 09 (Mezger et a l . 1974). On the basis of the Auer and Mihalas atmospheres which were generally accepted before 1978, i t is therefore safe to set R = I for the majority of giant HII regions. According to eq. (3.4), y+ measurements would then y i e l d the t o t a l He abundance d i r e c t l y . 4.

Previous Systematic Surveys The He recombination l i n e work up to 1973 which has been done mainly at 5 GHz with the

140- and 210-ft telescopes is summarized in paper I . The main results are as follows: (i)

y+~O.lO throughout the Galaxy

(ii)

HII regions close to the galactic center have systematically lower y+ values than spiral arm HII regions IONIZEDHEUUMAS A FUNCTIONOF INFERREDLYMANCONTINUUMPHOTONFLUX 0.14

f30.5 I 39.5 '

0'9

0 8'

BO II

09.5 I

07

o~

09 I

0 8 0 / 0 6 05 O L . ~ S U P E R G I A N T S II I I I I 05 '

04

~

ZAMS O

O

0.12 0 •

I0.10

• 0

0 0

0

•

0

•

0 •

÷

~0.08 tO.06 V

00

0

0

0.04 0.02 T

T

T T

10~

I

I

I

10~9

105o

1051

1052

Nclsec-') •

+

•

°

He abundances as a functzon of stellar Lyman contznuum photon flux N c' as inferred from radio continuum measurements. The corresponding spectral types are also shown (paper I).

Helium in the Galaxy

(iii)

363

the data suggest an inverse c o r r e l a t i o n between y+ and the Infrared Excess (defined below).

In Fig. 6 the y+ measurements from the 39 g a l a c t i c HII regions investigated are plotted against Nc', the number of s t e l l a r Lyman continuum photons required f o r the observed radio continuum f l u x density. These data are compatible with R = I over a large range of Nc' (and thus s t e l l a r e f f e c t i v e temperatures), as suggested by the model calculations (section 3) and, therefore, with y = 0.I0 throughout the Galaxy. There are a number of sources, however, which appear not to agree with this i n t e r p r e t a t i o n , notably those which have upper l i m i t s of ~0.02. Except f o r Ori B (leftmost l i m i t ) whose low He i o n i z a t i o n is r e a d i l y understood as a consequence of the low e f f e c t i v e temperature of i t s e x c i t i n g stars (Frey et a l . 1979), these low y+ sources are located at the Galactic center. Their continuum f l u x requires N ' >2.1050 photons sec - I the equivalent of several 06 stars. Since such stars have c y =0.40 according to the Auer and Mihalas model atmospheres, t h e i r surrounding HII regions should have R = 1 (Fig. 5). This contradicts the low y+ values observed i f y = 0.I0 is maintained. As an explanation i t was suggested in paper I that dust in HII regions is absorbing Lyman continuum r a d i a t i o n in such a way that He-ionizing photons (228
CORREL AT ION

F HE ABUNDANCE

N ITM I R EXCESS o [f~l'~lC..Jl~NqN~;

U~UILI~

tions (paper I ) .

AND 1,100RUQOO(19731

LIR IRE - - I Nc

UCL ~ T R

hu

The c o r r e l a t i o n may be r e a d i l y understood in terms of the dust optical depth T in the HII re-

g

ii

W

gion. On the one hand, an increasing T reduces the nebular Ly~ radiat i o n , thus increasing

Q

IRE. On the other hand, an increasing T causes y+ to decrease i f the

~ .

.

.

.

.

.

.

.

.

m .

.

.

.

.

.

.

.

.

I NFRRREO

.

1'0

EXCES<3

'

'

.

.

.

.

.

.

.

1B

I

Observed relation bundance to Infrared Excess (~merson and Jennings 1978).

C. Thum

364

dust is s e l e c t i v e l y absorbing. Fig. 7 shows y+ versus IRE as measured by Emerson and Jennings (1978) which shows the expected e f f e c t .

I t is noted, however, that an even better explanation

of t h i s c o r r e l a t i o n can be found without r e l y i n g on s e l e c t i v e absorption (Panagia 1980). A f t e r 1973 when the Effelsberg lO0-m telescope became f u l l y operational high q u a l i t y I09~ (5 GHz) and 91~ (8.6 GHz) He recombination l i n e measurements became available warranting a new analysis s i m i l a r to that in paper I . The main r e s u l t of t h i s work (Churchwell et a l . 1978; hereafter referred to as paper I I ) is that the v a r i a t i o n of y+ with g a l a c t r o c e n t r i c distance DG already indicated in paper I could be demonstrated (Fig. 8). The y+ d i s t r i b u t i o n was found to peak at DG

~9 kpc. The f a l l - o f f

from the peak value y+ =0.08 toward the outer

parts of the Galaxy was interpreted as a genuine decrease of y, whereas the decrease toward the g a l a c t i c center was a t t r i b u t e d to a decreasing R. This view of the y

+

v a r i a t i o n outside the solar c i r c l e is supported by the absence of

an observable geometric e f f e c t i f the 91a ( t e l e s c o p e r e s o l u t i o n 0 A = 1.5 arc min) and the 109a data (@A = 2.6 arc min) are compared. This means R ~1, and y+ ~y (Fig. 4 and eq. 3.4) w i t h i n observational e r r o r . A h i n t to why R should decrease towards the Galactic center came from the systematic v a r i a t i o n of HII region electron temperatures Te with DG as claimed in paper I I .

I t appears

now that despite the importance of selection effects (Brown et a l . 1978), the Te v a r i a t i o n can be interpreted as a v a r i a t i o n of the abundance of the coolants in the HII region, mainly oxygen. The lower Te towards the Galactic l

i

l

f

r

center would then r e f l e c t a higher heavy element abundance produced by the increased

o, Iz

stellar activity.

Since dust is made up p r i -

oJo •

•

" Q

"o

marily of heavier elements, s e l e c t i v e l y absorbing dust as postulate# in paper I may ex-

o.oB

p l a i n the observed behavior of R. We shall see

oo6 oo

.

in the next section, however, that selective

.

004

dust absorption can only play a minor role in HII regions and that another hypothesis based

0,02

on v a r i a t i o n s of the s t e l l a r UV r a d i a t i o n o o(

I

I

I

I

t

I

I

spectrum with the heavy element abundance can

b

account more r e a l i s t i c a l l y

oJo

f o r the observed

trend of R.

o o~

H; 7

5.

o.0~

New Model Atmospheres The observations presented in paper I and

0.0'~

I I suggested that the He+ volume is smaller

0.02

than the H+ volume f o r many g a l a c t i c HII re0.0(

o

~

~

~

~

,'o

,'2

,"

Do(Wpc)

gions, e s p e c i a l l y toward the Galactic center. The presence of neutral helium should manifest itself

Fi~. 8. Variation of the He + abundance (paper II). (a): results obtained for individual HII regions; sizes of dots refer to the accuracy of the measurements. (b): average value obtained from the indicated number of the individual sources.

p r i m a r i l y in the geometric e f f e c t as

discussed in section 3 and i l l u s t r a t e d in Fig. 4. The required high spatial r e s o l u t i o n data were taken with the lO0-m telescope in the l a s t 3 years. Thum et a l . (1980; hereafter referred to as paper I I I )

observed the 76~

Helium in the Galaxy

365

(14.7 GHz) and 66~ (22.4 GHz) He and H lines. At these high frequencies opacity and non-LTE effects

yl

I

I

I

as well as pressure broadening are negligible, ex-

0.08

cept maybe for the most compact HII regions. The

'w3A

*

0.04

results should therefore be d i r e c t l y interpretable in terms of y+. The telescope resolutions were 0.9

0.08

arc min and 0.7 arc min at 14.7 and 22.4 GHz, re-

0.04

spectively.

NGC 7538 Ori A -

0.10

In Fig. 9 these data are compared with lower

0.08

resolution observations obtained with the lO0-m 0.06

OR 21-

telescope. I t is seen that sources inside the solar c i r c l e (those below Ori A) clearly show a geometric

0.02

+

0.I0

effect, while in the two outermost regions the vari-

w4g

ation of y+, i f present at a l l , is less than 0.01. 0.06

These data independently confirm the suggestion of

w51

0.10 0.06

+++

0.10

papers I and I I that the decrease of y+ towards the Galactic center is due to a shrinkage of the He+

+w,3

volume relative to the H+ volume. The presence of neutral He in HII regions, increasing in proportion

0.06

~,

0.06

towards the Galactic center, is therefore highly

w 31

probable.

4, 0.02

0.08

This is clearly at variance with the notion

Sgr B2

that the excitation of giant HII regions is mainly

0.04 [

i

I

I

2

3

provided by 06 stars (section 3) with an effective y value of 0.40 (see below). I f selective absorption

TELESC0PE HPBW[arcmin]

by dust were to be a way out of t h i s discrepancy, i t would have to reduce the s t e l l a r y value by more than a factor of 2, namely below y =0.20 required ~ i

Variation of the He + abunth telescope half power width. Sources are arranged (from top to bottom) in order of decreasing galactocentric distance: 12.4, 11.8, 10.4, 10.0, 9.4, 7.8, 5.4, 5.1, 0.1 kpc. From paper III.

y. I

'

l

'

,

I I00

i

~

'

I

i

~

'

for f u l l He ionization. This would require the ratio ao of the dust absorption cross section for He and H ionizing photons to be ao = 4 ± 1 (Panagia and

i

~RIONIA

'

I

l

i

I

0.10F

oo40.0t-

117 [ , ~

200

300 r("}

400

I 580

I-

600

~ 0 . Observed variation of HeY--abundances in the Orion Nebula as a function of the distance from the principal ionizing star gIC Ori (Pankonin et al. 1980). Circles refer to 109~ measurements, crosses to optical data (Peimbert and Torres-Peimbert, 1977). The two solid lines are linear fits to radio data taken at positions east and southeast of @IC Ori (lower line) and west of eiC ori (upper line).

366

C. Thum

Smith 1978). Such a high value postulates a steeper than linear increase of the dust absorption over the wavelength range 228 - 912 ~. The fact that no known dust material has this property (Huffman 1977) led to a search for alternative solutions. Several other i n v e s t i g a t i o n s came to conclusions incompatible with s e l e c t i v e absorption. Pankonin et a l . (1980) reported He and HI09~ measurements from a number of positions in the Orion Nebula. The y+ values obtained are presented in Fig. I0 as a f u n c t i o n of the radial distance r from the p r i n c i p a l e x c i t i n g star 01C Ori (spectral type 06). I t is evident that y

+

decreases systematically from i t s peak value at r = O, demonstrating that R < i . I f s e l e c t i v e dust absorption were responsible for the decrease, a c o r r e l a t i o n of y+ with the e x t i n c t i o n at optical wavelengths would be expected. Such a c o r r e l a t i o n is not observed (Pankonin et a l . 1980). Evidence against the Auer and Mihalas atmospheres and thus against the necessity of invoking selective absorption came from an optical investigation of the large c i r c u l a r HII region around the 08 star XOri (Harms et al. 1978). Observing the forbidden X5007 line of 0++ at several locations in the nebula, the authors found a substantially smaller 0++ volume than expected from a non-LTE model atmosphere of an 08 star. In this nebula selective absorption can certainly not resolve the problem since there is not enough dust. The optical extinction of AV = 0.39 towards XOri implies a dust UV absorption optical depth of less than 0.1 mag. The authors conclude that at wavelengths shorter than 353 R, the ionization edge of 0++, the star emits much less photons than predicted by the non-LTE atmospheres. A similar suggestion was made by Panagia and Smith (1978) in their interpretation of the y+-IRE correlation. I t was suspected that the Auer and Mihalas atmospheres were not correct, since they neglect line blanketing in the far UV. In fact, new model atmospheres which include l i n e blanketing (Kurucz 1979) predict a much steeper UV spectrum due to metal absorption lines crowding at X<500 R. Hence, the corresponding y values are considerably reduced. Fig. 11 shows y as a function of the s t e l l a r effective temperature or spectral type for the new model at-

09 08 07 08.5 07.5 06.5

Sp.T. 09.5 0.7

0.6 0.5

:ff

0.4 -

'

'

'

0"03

'

'

'

(n} BLACKBODY (b) AUERANDMIHALAS(1972) (c} KURUCZ(1979} (c') LOGg=50 (c"} LOGg=4.5 (c"') LOGg=4.0

,=, 0.3-

0.1

'

06

05.5

05

04

'

'

'

'

'

I-

-

(b)

o

.I oI J"

"

~

~,(c') ~(c")

//"

// / (c,,,)//i// _1-

03

. - , ..(o)R=O95for y

-

3.5

I

I

4

4.5

I

5

I

5.5

Teff/104K F_ig. 11. Ratio X of He and H ionizing photons emitted by 0 stars v~. spectral types (Sp.T) and effective temperature Teff. Curves a, b and c' refer to a blackbody and 2 different model atmospheres. The X values at which He is nearly fully ionized (R = 0.95) for 3 different He abundances are indicated by arrows. From Mezger (1980).

Helium in the Galaxy

367

mospheres together with the Auer and Mihalas and blackbody atmospheres. The new y Values are more than a factor of two below the Auer and Mihalas curve. The 06 star e x c i t i n g the Orion Nebula would thus be barely able to f u l l y ionize He. This is in good agreement with observations (Fig. i0) and thus lends confidence in the new atmospheres. Since the e f f e c t i v e y value of an OB star c l u s t e r is close to that of a single 06 s t a r , we now expect that He is only marginally f u l l y ionized in giant HII regions, R ~ I . Under these conditions small changes of the average e f f e c t i v e temperature of an OB star c l u s t e r d r a s t i c a l l y a l t e r the He i o n i z a t i o n structure (Panagia 1980). Such changes can be caused by v a r i a tions in Z, the heavy element abundance, as envisaged in paper I I (section 4). A s l i g h t rise of Z above i t s solar value w i l l increase the X<500 ~ opacity in the 0 s t a r ' s atmosphere and therefore reduce y . Since R ~I f o r solar abundances already, such a picture would y i e l d R s u b s t a n t i a l l y smaller than u n i t y , and thus an e a s i l y observable e f f e c t . This idea w i l l be discussed in the next section. I t thus appears that the UV spectra derived from the new model atmospheres can explain the observed behavior of R. In view of the d i f f i c u l t i e s

described above, we conclude that

selective dust absorption plays only a minor role in real HII regions, the dominant e f f e c t being the v a r i a t i o n of e f f e c t i v e temperatures of the e x c i t i n g stars. 6.

The Variation of the Helium Abundance across the Galaxy In the presence of a substantial f r a c t i o n of neutral helium, the best estimate of the

t o t a l helium abundance is obtained from high spatial r e s o l u t i o n / h i g h frequency observations as a r e s u l t of the geometric e f f e c t and related considerations (section 2). The 66~ and 76~ data (paper I I I )

therefore represent the best sample of He abundances c u r r e n t l y available f o r

studies of the Galactic He v a r i a t i o n . I

I

I

I

I

I

I

I

I

I

I

I

I

I

I

I

0.12

0.10

L~

O.Oe y* 006

e~

OOh

002

I

0

I

I

2

I

I

4

I

I

6

I

I

8 0 G [kpc]

I

I

10

I

I

12

I

I

I

14

Ionized helium abundance y+ as a function of galactocentric distance D G (paper III). Boxes, shaded areas and dashed lines refer to 109a data (paper II) as described in the text. Individual sources which have y+ determined at several transitions are showh as dots connected by vertical lines.

368

C. Thum

Fig. 12 shows these data ( f i l l e d dots l a b e l l e d 66 and 76) together with a rearrangement of the 109m data. The sources are sorted into Galactic annuli of width ADG which has been chosen so that each annulus contains an equal number (5) of 109m sources. The range of 109m y+ values + values

w i t h i n each annulus is o u t l i n e d by the boundaries of the boxes, while the median y

are indicated by dashed l i n e s . In an attempt to reduce the e f f e c t of possible unnoticed experimental errors the extreme y + y values in Fig. 12.

+

values of each box are omitted leaving the shaded range of

Besides the trends already noted in section 4 i t is clear from this figure that y+ shows a spread of values which widens towards the Galactic center. This spread appears to be much too large to be f u l l y accounted for by real y variations. Variations of Ro, the quantity relevant for observations, within a galactic annulus are a much more plausible explanation. I f R is decreasing steadily the range of possible values R widens, accounting naturally for the +

0

observed spread of y . Sources f o r which 66m and 76m data are a v a i l a b l e are superimposed on Fig. 12, together with t h e i r 91/109m data, as f i l l e d dots connected by v e r t i c a l l i n e s . Inside the solar c i r c l e most 66a points f a l l at or above the upper boundary of the 109m d i s t r i b u t i o n . This behavior, again, finds an explanation in the geometric e f f e c t , since both the 66m data and the 109m upper boundary tend to r e f e r to resolved regions where y

+

.

is expected to approach i t s maximum

value, y. A combination of these 66m points and the 109m upper boundary can therefore be regarded as a good lower l i m i t to the He abundance in the i n t e r s t e l l a r medium at 7 ~DG ~15 kpc.

In Fig. 13, this lower l i m i t is combined with optical He abundance determinations, mainly for planetary nebulae (Torres-Peimbert and Peimbert 1977). In paper I l l i t is argued that y from planetary nebulae, PN, represents an upper l i m i t to the He abundance in the i n t e r s t e l l a r medium, y(ISM) ~y(PN). This is certainly correct in the solar neighborhood, and can be regarded as a good working hypothesis for the whole distance range considered here. I t implies that the enrichment of PN in He from nucleosynI

I

I

I

1

•

I

I

I

I

0.15 y

thesis in the parent star is at least equal to the He enrichment of the ISM during the l i f e -

(y')

0]4

time of the star. The shaded area in Fig. 13 •

0 °

should therefore represent the He variation in

0.13

the ISM with reasonable accuracy. I t is well 0.12

approximated as

0.11 Alog Y/ADG = -0.015 kpc

-1

0.10

where Y now refers to the mass fraction.

OOg 0.06

Primordial Helium

007

Helium in the Galaxy is thought to be sum of primordial He, Yp, synthesized in the Big ?

8

9

I0

l]

IZ

13

OG[kpc]

]4

15

Bang, and a term AY, describing the He synthe-

sized in stars and subsequently returned to the Helium abundances y from optical lanetary Nebulae (dots and crosses) ISM. and HII regions (squares). He + abundances from radio HII regions (symbols as in Fig.12). Y = y +Ay The shaded area indicates the probable range P of y in the interstellar medium. From paper III.

~

p

Helium in the Galaxy

369

Yp can best be estimated from Y when the term AY is smallest. This is the case at the edge of the Galaxy where s t e l l a r He production is low due to the l e v e l l i n g - o f f of the s t e l l a r b i r t h rate (Smith et a l . 1978). From Fig. 13 Yedge = 0.08 ±.01 is i n f e r r e d , corresponding to Yedge = 0.24 ± .02. However, yp must be lower than t h i s , because even at large DG the observed y is contaminated by helium produced in stars. The very existence of HII regions at DG =15 kpc proves the presence of young massive stars which are the main contributors to the He enrichment of the ISM. Other evidence comes from inferences about of the i n t e r s t e l l a r abundance of oxygen (Mezger et a l . 1979) which is produced in even more massive stars. These data show that O/H or, more generally, the heavy element abundance Z at large DG is not very much reduced over i t s solar value. Applying the r e l a t i o n AY/AZ ~2 u s u a l l y adopted f o r the enrichment of the ISM with He and heavier elements leads to a s t e l l a r contamination of Yedge of up to ~Z).02. The estimate of the primordial value is therefore yp = 0.07 ± .01

or

Yp = 0.22 ± .02

(6.1)

I t is informative to compare t h i s r e s u l t , which agrees well with a recent optical determination from gas-rich external galaxies (Lequeux et al. 1979), with predictions from nucleosynthesis calculations for the Big Bang (paper I I ) .

According to the standard theory

(Wagoner 1973) the observed value Yp (eq. 6.1) corresponds to a present nucleon density of ~I0 -31 g cm-3.

Since the matter in v i s i b l e galaxies alone can account f o r t h i s density, sub-

s t a n t i a l c o n t r i b u t i o n s to the g r a v i t a t i n g mass in the Universe from unseen matter appear to be excluded in the framework of t h i s standard theory. The Universe would then be open and expanding forever. I t has recently been suggested that new kinds of neutrinos may have influenced the evol u t i o n of the early Universe. Calculations by Yang et a l . (1979) predict Yp = 0.250 f o r the case of three neutrinos (the f a m i l i a r e- and ~-neutrinos and a possible T neutrino) and the present mean density (2.10 -31 g cm-3). While t h i s r e s u l t may s t i l l

be compatible with obser-

vations, a higher number of new l i g h t neutrinos appears h i g h l y u n l i k e l y (Schmid-Burgk 1980). Helium in the inner Galaxy In the inner parts of the Galactic disk, DG <8 kpc, y cannot be traced by optical means. There, the radio data in Fig. 12 represent the only information. Again, a r e l i a b l e lower l i m i t to y (ISM) may be obtained from the high frequency data or the upper boundary of the i09~ d i s t r i b u t i o n , whichever is higher, y (ISM)

y

+

~0.I05 is estimated f o r DG <8 kpc.

Recent observations at 15 GHz with the lO0-m telescope ( A l t e n h o f f et a l . 1981) yielded values for 15 more HII regions inside 8 kpc. These new high r e s o l u t i o n data continue the

trend towards higher He+ abundances as observed in paper I I I

(Fig. 9) and as expected from

the geometric e f f e c t . Values of y+ ~).12 are measured for several sources including one close to the Galactic center (G 0 . 5 - 0 . 1 S ) . This indicates the lower l i m i t to y has to be revised upwards f o r DG <5 kpc. The lower boundary of the i09~ y+ data in Fig. 12 reaches values below 0.02 at the Galact i c center. Since at that distance the data refer to unresolved regions, the q u a n t i t y R can <

be estimated as ~0.2. I f the non-LTE atmospheres are adopted and selective absorption is d i s carded, t h i s would correspond to y ~0.03 or e x c i t a t i o n by early B stars (Fig. I i ) as was suggested by Rodriguez and Chaisson (1979). Apart from the question of the model atmospheres (section 5), t h i s suggestion faces serious problems. To ionize a t y p i c a l giant HII region, a

370

C. Thum

huge number of B stars would be required. To avoid the 0 stars normally accompanying them in a large c l u s t e r , one would have to assume that they a l l have an age of ~I07 yrs, so that the B stars would s t i l l

be on the main sequence while the 0 stars would have evolved o f f i t .

Since

there are quite a few giant HII regions in the Galactic center region, a l l of which have R ~0.2, this would require that they are a l l excited by B stars of nearly the same age, an u n l i k e l y s i t u a t i o n . A more d i r e c t argument bears on the observation of H20 maser sources, mainly in Sgr B2, which indicate the presence of young massive stars (Genzel and Downes 1977). Instead of the B star hypothesis another explanation is favored which follows n a t u r a l l y from the new model atmospheres and a systematic increase of the heavy element abundance Z towards the Galactic center. We noted in section 5 that an increased Z w i l l r e s u l t in a lower ¥ value because of increased opacity at ~<500 ~. Panagia (1980) suggests two more consequences of an increase in Z, both having the e f f e c t of decreasing ¥. Since higher opacity will

increase the radius of a star, i t s e f f e c t i v e temperature w i l l decrease. The other e f f e c t

is a reduction of the upper c u t - o f f of the i n i t i a l

mass function due to higher Z. Since at

higher Z the cooling rate of the i n t e r s t e l l a r gas r i s e s , clouds of smaller masses can become g r a v i t a t i o n a l l y unstable and form less massive stars. In a d d i t i o n , the opacity of the mater i a l accreted onto a forming star increases thus reducing i t s f i n a l mass. With fewer high mass stars in a c l u s t e r , i t s e f f e c t i v e y w i l l decrease. Panagia (1980) estimates that the upper c u t - o f f of the I n i t i a l Mass Function is more important than the effects of blanketing by UV l i n e s and s t e l l a r r a d i i , although none of these is n e g l i g i b l e . Observational evidence f o r a large scale gradient of the heavy element abundance over the Galactic disk has been obtained from HII region electron temperatures (paper I I ; Mezger et a l . 1979). Other evidence, mainly from o p t ic al HII regions and planetary nebulae f o r a smaller range of galactocentric distance, has been summarized by Panagia (1979). Based on these i n v e s t i g a t i o n s , Z in the inner parts of the Galaxy appears to be higher by a f a c t o r 2-3 than in the solar neighborhood. Assuming a Z gradient of this order and estimating the Z-dependence of the above 3 effects on the He

0.15

I

I

I

I

i o n i z a t i o n Panagia (1980) can near!y reproduce the observed trends of y+, f o r both unresolved

Y ~ (y÷)

TOTAL

and resolved HII regions (Fig. 14). He finds a decrease of the average e f f e c t i v e tempera-

,

0.I0

ture of the 0 stars from 39 000 K at the sun to 31 000 K at the Galactic center, and infers a t o t a l He abundance y >0.12 inside 5 kpc. Al-

•

/

0.05

t

HELIU~

though the best f i t

shown in Fig. 14 w i l l need

some r e v i s i o n in view of the new data from the

/.~

lO0-m telescope, i t seems from the f i g u r e that these model c a l c u l a t i o n s can s a t i s f a c t o r i l y account f o r the observed behavior of the Fie

0

5

10

15

i o n i z a t i o n in the Galaxy.

gG[kpc] 7.

~

e

The variation of the total He + of unresolved HII-Regions (thin line) and of fully resolved HII-Regions (dashed dotted line) after model calculations of Panagia (1980). Crosses refer to He + measurements of (unresolved) giant HII regions (Fig. 8).

Conclusion The work presented in this paper spans

more than one and a h a l f decades of research with the world's largest telescopes. The main goal of this research, namely the determination of He abundances throughout the Galaxy, has

Helium in the Galaxy

371

been reached to a considerable extent as shown in Figs. 12-14. This progress was possible because the physical processes involved in the He and H recombination l i n e transfer now appear to be understood. I t has been shown by several investigators that He+ abundances derived under the assumption of LTE are r e l i a b l e lower limits to the total He abundance (section 2). These investigations furthermore suggest that observations at higher frequencies (typically above 5 GHz) yield better values, mainly because of higher resolution, but also because the disturbing effects of opacity, pressure broadening and stimulated emission tend to be negligible there. Another important step forward has been possible on the basis of new model O star atmospheres which include blanketin0 by far-UV lines, namely of heavy elements. The often-observed non-coincidence of He+ and H+ volumes may thus be seen to be a consequence of the ionizing spectrum of the exciting stars (section 5). I t remains to be seen, however, to what extent the new atmospheres are characteristic of real 0 stars with rotation and mass loss. In any case, the new atmospheres may allow a linking of two basic observations, namely the increase of the heavy element abundance Z towards the Galactic center and the shrinking of the He+ volume r e l a t i v e to the H+ volume (section 6). In a quantitative analysis of the effects thought to be associated with a Z gradient, Panagia (1980) obtained a reasonable f i t to existing He observations. This suggests that many important effects influencing the He+ abundance are presently understood. Despite this progress, two central questions are s t i l l not answered s a t i s f a c t o r i l y . F i r s t l y , what is the run of the He abundance in the inner Galaxy, and in particular, what is y at the Galactic center? For the range 0 ~ DG ~ 7 kpc in which most of the enrichment of the ISM by s t e l l a r nucleosynthesis products occurs, we have only a lower l i m i t on y. Secondly, what is the accurate value of the primordial He abundance? The discussion in the preceding section showed that higher accuracy is highly desirable i f measurements are to be compared with Big Bang models. Both questions demand higher precision of the measurements of the He abundance than have been available up to now. On the observational side, an increase in sensitivity by a factor of ~2 can be expected over the next 5 years. Thus, more sources w i l l be observable in the inner parts of the Galaxy, hopefully improving the He+ determination in a s t a t i s t i c a l sense. Another l i m i t a t i o n , especially for the distant sources in the inner Galaxy, is telescope resolution. Substantial progress in this respect w i l l not be possible before sensitive interferometers operating at high enough frequencies can eventually be used for spectroscopy. Unfortunately, the most interesting sources, i . e . those closest to the Galactic center, w i l l on average have the smallest r e l a t i v e He+ volume, thus requiring the largest correction for neutral helium. This unfavorable situation w i l l only change i f yO can be inferred in a more direct manner from observations of other ions with similar ionization potential. A good candidate for this approach seems to be sulfur, since its most abundant ions (in HII regions), S I l l and SIV, have infrared transitions and are thus in principle observable over a large galactic volume. ~ith the help of HII region models which relate the degrees of ionization of the relevant elements to the s t e l l a r spectrum, one may hope to derive ionization correction factors for the unseen neutral helium. Finally, i t may be possible to learn more about the young massive stars and their interaction with the ISM, especially those with masses around 4 M8 which contribute most to the He enrichment of the i n t e r s t e l l a r medium. Such elaborate corrections for yO may not be needed for the HII regions relevant for the determination of yp, since their neutral helium fraction is not expected to be substantial.

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The main observational problems here are the scarcity of sources at the edge of the Galaxy and the generally low surface brightness of these sources. Optical spectroscopy may be able to contribute here since these HII regions p r e f e r e n t i a l l y l i e in regions of the Galaxy which are less obscured by dust. But, without doubt, both branches of astronomy w i l l have to work hard i f the determination of the fundamental cosmological quantity Yp is to be substantially improved in the future. Acknowledgement I wish to express my gratitude to P. G. Mezger and L. F. Smith who introduced me to helium recombination line astronomy. I also want to thank P. G. Mezger, N. Panagia, J. Schmid-Burgk and J. Wink for c r i t i c a l l y reading the manuscript and Ann Downes for correcting the English text. References Altenhoff, ~. J., Martin-Pintado, J., Mezger, P. G., Thum, C., Wink, J. Astrophys., to be published Auer, L. H., Mihalas, D. Batchelor, A. S. J. Brocklehurst, M.

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