DETERMINing the

Metallicity in the

Galactic Center

Rachel L. Smith Examination Proposition II

October 19, 2007

MIRLIN (Mid-Infrared Large-well Imager) image of the Galactic center at 9, 13 and 21 µm (Image credit: Morris et al., Keck II Telescope).

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Determining the Metallicity of the Galactic Center

MIRLIN cover image with the most notable objects labeled. (http://irastro.jpl.nasa.gov/GalCen/galcen.html)

3

I. Abstract

We propose to use far-infrared emission lines of accessible ionization stages of H,

N, O, Ne and Ar from five H II regions within the inner 100 parsecs of the Galactic

center to determine abundance ratios and reliable metallicity values for the galactic

center, and to explore the possibility of a metallicity gradient near the Galactic nucleus.

An accurate determination of the metallicity of Galactic center gas bears upon several

critical aspects of galactic evolution, including star formation in the Galactic center, the

heating and cooling processes for Galactic center clouds, estimates of cloud masses

inferred from molecular line observations, and molecular cloud chemistry.

II. Galactic center metallicity

Introduction: The Galactic center is the closest of all galactic nuclei in the

Universe, and one which shows enhanced star formation compared with other parts of the

Galaxy. The Galactic center is characterized by several unusual features, including a

supermassive black hole, (Sgr A), the densest star cluster in the Galaxy, a central H II

region (Sgr A West), a torus of circumstellar gas (White et al. 2007, in preparation), and

the Galaxy’s most massive interstellar clouds. Stellar birth and death is quantitatively

higher in the Galactic center in comparison to other regions in the Galaxy, and has

supernova explosion and planetary nebula formation rates that are each roughly two

orders of magnitude greater than those in the Solar neighborhood (~ 8 kpc from the

Galactic center).

Previous studies predict that star formation within 10-100 parsecs of the Galactic

nucleus is strongly affected by the physical extremes of this region; these are strong tidal

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forces, large internal velocity dispersions and large magnetic fields. These characteristics

are predicted to favor the formation of massive stars, such as the spectacular arches

cluster shown in Figure 1, thereby enhancing the metallicity in the central galactic

regions (Morris 1993; Morris 2005).

Importance of Galactic center metallicity: The high efficiency of star formation in

the Galactic center and observations consistent with an enhanced “nuclear maturity” of

Galactic center gas compared to the Solar neighborhood presages an enhanced metallicity

due to nuclear processing, evident in increases in primary (H-burning on carbon

generated within the star) and secondary (H-burning by CNO processing on originally-

present carbon) nuclear burning products. Observations of enhanced stellar abundances

of 13C/12C, 17O/16O and indications for increased N/O in the galactic center are consistent

with metal enrichment (Wannier 1989 review) and one expects the metallicity of Galactic

center gas to be enhanced (Oort 1977). However, to date there are no direct and reliable

measures of this metallicity.

The study of chemical abundances and their variation within the Galaxy (as well

as from one galaxy to another) is of fundamental importance for our understanding of

Galactic evolution in general (Shaver et al. 1983), which in turn is critical toward

understanding solar system origins. An exact knowledge of metallicity in the Galactic

center, as well as of a potential gradient within the inner 100 pc, are important for

constraining modeling parameters of Galactic evolution (Gusten & Ungerechts 1985) and

for understanding Galactic cloud chemistry, and heating and cooling processes (Morris et

al. 1983; Guesten et al. 1985; Morris 1993).

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The role of metallicty on star formation is not precisely understood. Enhanced

metallicity, expected for the Galactic center, affects the initial mass function (IMF)

inasmuch as it determines the opacity of the star-forming medium. The opacity, in turn,

determines the rate of energy loss (and thus the rate of contraction) in a contracting

molecular cloud. The greater the opacity, the more time a cloud core has to contract and

gather material before accretion ceases; all of these effects favor the formation of massive

stars (Morris 1993). A more thorough understanding of metallicity in the complex

environment of the Galactic center thus bears upon a more comprehensive understanding

of star formation in general and evolution of our solar system in particular.

H II regions: H II regions provide the most accessible probe of current interstellar

abundances of the heavy elements. They yield a fossil record of the nucleosynthetic

enrichment that has taken place in successive stellar generations, and enable a tracing of

this evolution via chemical history within and between galactic systems (Shaver et al.

1983).

H II regions are diffuse clouds of gas in the interstellar medium surrounding

massive O- and B-type stars. Ultraviolet radiation from these stars ionizes hydrogen and

other atoms within a volume of space that defines the region of ionization, or “Strömgren

sphere.” Recombination lines within these areas of high ionization are responsible for the

colors associated with many H II regions. Unfamiliar spectral lines (the so-called

forbidden lines) of familiar elements such as oxygen, nitrogen, neon and argon, are

produced by collisional excitation of atoms or ions of these elements from ground

electronic configurations to nearby levels where they cascade back with the emission of

radiation (Aller 1987).

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H II regions are interspersed throughout our Galaxy; Figure 2 shows many H II

regions (red patches) interspersed throughout the spiral arms of the nearby Whirlpool

galaxy (M 51). H II regions range from a few to several hundred parsecs in diameter

surrounding massive hot, blue OB stars, and have electron temperatures that have been

shown to vary with Galactocentric distance (Shaver et al. 1983; Afflerbach et al. 1997,

1997; Deharveng et al. 2000) from approximately 5,000 K in the Galactic center to

~10,000 K at 15 kpc, which is important in determining accurate effective temperatures

for abundance ratio calculations (Shaver et al. 1983).

For embedded, high-excitation H II regions, observing in the far-IR is preferred

over optical wavelengths for several reasons, including the insensitivity to extinction by

dust, located either within the ionized gas or in the neutral foreground material, of far-IR

as compared to optical wavelengths, thus leading to smaller ionization correction factors

and lower uncertainties in the derived line ratios (Rudolph et al. 1997). H II regions in the

Galactic center suffer from 30 magnitudes of visual extinction, and so observations of

these objects depend on far-IR techniques.

III. Investigating H II regions

Infrared Space Observatory: Operated under the European Space Agency (ESA) ,

the Infrared Space Observatory (ISO) was the world’s first true orbiting infrared

observatory. It’s operational phase was from 1995 to 1998, and it made ~ 30,000

individual imaging, photometric, spectroscopic and polarimetric observations from the

solar system to extra-Galactic sources. It was equipped with two spectrometers probing

long- and short-wavelengths, (LWS and SWS, respectively), as well as a camera

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(ISOCAM) and an imaging photo-polarimeter (ISOPHOT), jointly covering wavelengths

from 2.5 to 240 µm with spatial resolutions ranging from 1.5 (at the shortest

wavelengths) to 90 arcseconds (at the longer wavelengths) (ISO website).

ISO was groundbreaking in that it was the first spectroscopic infrared satellite

permitting far-infrared measurements of fine-structure lines to estimate the abundance

ratios N/H, N/O, Ne/H and Ar/H.

Atomic fine-structure lines: Infrared lines serve as valuable probes of accessible

ionization states within H II regions (Herter 1989). Tables 1 and 2 list the accessible

ionization states using the Infrared Space Observatory’s Long-Wavelength and Short-

Wave Spectrometers, respectively, collected for this study. These lines represent

transitions within various electron configurations of the respective ions. For example, the

[Ne II] line at 12.8 µm corresponds to the transition in the ground configuration of singly

ionized Neon. Neon has the ground electron configuration 1s22s22p5, where the letters s,

p, d indicate electrons with orbital angular momentum quantum numbers, l = 0, 1, 2. The

superscript denotes electrons in each shell. In the unfilled 2p shell the resultant electron

spin (S) and total electron orbital angular momentum (L) quantum numbers interact and

cause a splitting of each state into 2S +1 levels. For Neon, the lowest energy state where

the splitting occurs is from total angular momentum of

!

12

to

!

32

, with an associated

energy loss of roughly 780 cm-1.

The term fine-structure lines denotes jumps between individual levels of a ground

configuration. Brackets denote these transitions as forbidden; they are forbidden because

there is no parity change (i.e. Δl =0) when outer electrons jump between spin-states

8

within the same p orbital, for example. These lines are detectable in very low-density

regions (~ 103 ions/cm3), such as in H II regions and planetary nebulae (Aller 1987).

The great astrophysical value of forbidden lines is that certain of their intensity

ratios yield important diagnostic information for the region being studied-- such as the

electron density and electron temperature—which in turn is used to deduce the various

molecular abundances in the gas (Aller 1987).

IV. Previous relevant studies

While a reliable determination of Galactic center metallicity is still lacking,

several previous studies have shown important trends in nuclear maturity and high

metallicity in the Galactic center. Far-IR observations have shown high isotopic

abundance ratios (i.e. 13C/12C and 18O/16O), albeit with large errors, compared to solar

values, in support of the prediction of increased nuclear processing in conjunction with

high rates of star formation in the Galactic center (Wannier 1989). Observations of the 7-

µm fine-structure line in Sgr A, the H II region in the center of the Galaxy, have shown a

factor-of-two enrichment in argon compared with the vicinity of the Sun (Willner et al.

1979). However, the uncertainty in this determination is large due to complexities in the

interpretation of broad emission lines with many velocity components (Willner et al.

1979), and potential underestimation of electron temperature and collision strength for

Ar+, both of which would lead to a spuriously high abundance (Lester et al. 1981). A

similar factor-of-two enhancement in Brα (4.05 µm), [Ne II] (12.8 µm) and [Ar III] (9.0

µm) was noted in emission from Sgr A West, but the metallicity was not calculated due

to the unknown population distribution of the ionization states (Lacy et al. 1989).

9

The abundance ratio of N/O is important in establishing the importance of

primary vs. secondary production of nitrogen as a function of metallicity. In high-

metallicity environments nitrogen is believed to be synthesized via the CNO cycle in

intermediate-mass stars, and thus nitrogen is considered a “secondary” element, with N/O

increasing linearly with O/H (Renzini & Voli 1981). However, evidence for primary

production via a constant N/O signature in low-metallicity environments have also been

observed (e.g. Garnett 1990; Thuan et al. 1995), indicating nitrogen production via a

primary process, such as by successive dredge-ups of enriched cores in intermediate-mass

stars (Renzini & Voli 1981).

However, interpreting N/O ratios is complicated by several factors, including the

enrichment of O vs. N in the Galactic center as a result of supernovae from a stellar

population dominated by high-mass stars, while in the outer regions of the disk

overproduction of N may result from CN processing in a lower-mass stellar population

where supernovae are less frequent. For example, far-IR observations of [N III] and [O

III] fine structure lines in H II regions has been used to infer an N/O ratio, and a factor of

two enhancement in N/O in the Galactic center have been found (Lester et al. 1987;

Rubin et al. 1988), with the assumption that the N++/O++ ratio is a reasonable

approximation of the molecular abundance. However, while results indicate an

enhancement in N/O in Sgr A and the 5 kpc ring, which could be explained by enhanced

star formation, the N/O ratio in H II region G0.5-0.0 is enhanced by a factor of 2 over Sgr

A and the disk at 5 kpc (Lester et al. 1987).

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V. Proposed study

We propose to use far-infrared line fluxes measured by Morris, van Dishoeck and

Habing in 1995 with ISO’s Short-Wavelength Spectrometer (SWS) and Long-

Wavelength Spectrometer (LWS). Observations were made using the Fabry-Perot high-

resolution interferometer in order to obtain the highest possible line/continuum ratio.

Ratios will be calculated using these data and existing radio continuum observations to

construct temperature, density and radiation field models within each H II region.

Studies of H II regions show a decrease in metallicity with distance from the

centers of our own (e.g. Shaver et al. 1983; Giveon et al. 2002; Martin-Hernandez et al.

2002) and other galaxies (i.e. Aller 1984; Dinerstein 1990). Therefore, in addition to

abundance ratios, we will also explore any dependence of Galactic center metallicity on

Galactocentric radius (RG) within the Galaxy’s inner 100 pc, which will be especially

interesting given the inference that the Galaxy has an abundance gradient for all but the

very central H II regions (RG ~ < 3 kpc), where the electron temperature Te (RG) is flat

(Wink et al. 1983).

Selected H II regions: Five H II regions were selected: G-0.02-0.07 (shock

region), G-0.02-0.07 (metallicity region), AFGL 5376-1, G0.18-0.04 and Sgr C. They

were selected based on their strong infrared emission and relative isolation from each

other within the Galactic center; all are within 75 pc of the Galactic nucleus in projection,

yet have a spread of Galactocentric distances from one another so that any sharp

metallicity gradients near the nucleus might be probed. They are also all relatively well

understood, with existing velocity, mapping and absorption studies indicating that they

are actually close to the Galactic center rather than coincidentally superimposed.

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Analytical method, overview: Tables 1 and 2 list the accessible ionization states

probed by LWS01 and SWS01, respectively, for sample H II metallicity region G-0.02-

0.07. Sample short-wavelength spectral lines are shown in Figures 3. All accessible and

reasonable ionization states of O, N, Ar and Ne were detected. As noble gases, Ar and Ne

are unlikely to be depleted onto grains and thus were chosen as probes for the metal

abundance. The H I Pf-α (7.45 µm) lines were included for direct comparison with argon

and to provide an estimate of the hydrogen column density for comparison with radio

continuum measurements and because the hydrogen abundance is fundamental to

knowing and expressing metallicity. HI Br-α (4.05 µm) is included as a constraint on the

mid-IR extinction by comparison with H I Pf-α (elaborated further below). Since [O II]

has no accessible transitions in the infrared, we intend to estimate N/O ratio from the [N

III]/[O III] intensity ratios, using models of H II regions and resulting predicted ionic

abundances to apply corrections to the resulting N/O ratio. Relative ionic abundances will

be calculated using the procedure outlined in the next section.

In order to make an exact determination of the abundance of an element, one

must have measured the total line flux of at least one line from each ionization state of

the atom. In practice, however, only a few ionization states contribute to the total

abundance, as a complete set of lines from the relevant ionization states are generally not

obtainable and various models are needed to correct for the unseen ionization states

(Rudolph et al. 1997).

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VI. Analytical method: steps

1. Correct far-IR fluxes for interstellar extinction: While smaller at infrared than optical

wavelengths, mid-IR extinction still exists and needs to be corrected for. In this study,

extinction is constrained via comparison of Br-α and Pf-α H I lines, which will be used

to derive the extinction using recombination theory, i.e. ratios of our lines will be used

against established models of the hydrogen recombination cascade created using

recombination lines collected at radio wavelengths (not subject to extinction). We will

compute a differential extinction factor by comparing our line fluxes and recombination

lines from theory to derive the extinction correction factor for each source.

2. Using line fluxes to determine abundance ratios

2a. Theory: Generally, we define the line emissivity (per unit volume, per unit

time) for an optically thin spectral line as,

!

" #ij( ) = N jA ji

hc

#

!

j > i( ) (Dwivedi and Gupta 1994)

where Aji is the spontaneous radiative transition probability and Nj the number density of

the upper level j. Emissivity is important for determining ionic abundance ratios between

any two ions, and is thus an important parameter in determining the final total atomic

abundances in a source. Alternatively, we can define the volume emissivity,

!

j"Ne,Te( )

for each line λ for a given ion X+i with density

!

NX

+i (cm−3), electron density Ne (cm−3)

and electron temperature, Te; this emissivity is propagated through the remainder of the

analytical procedure. The relation for normalized volume emissivity,

!

" Ne,Te( ) , is,

!

" Ne,Te( ) =j#

NX+iNe

(Simpson et al. 1995),

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and the flux in any optically thin line corrected for extinction is then given by,

!

F" = #"NX

+i NedV 4$% (Simpson et al. 1995),

where

!

dV = dld", the volume element of the beam,

!

" (units of solid angle), of the

telescope. Assuming a constant

!

"# over the source, the abundance ratio for any two ions

is equal to the ratio of the observed fluxes divided by the ratio of the appropriate

!

"#’s

(Simpson et al. 1995).

The ionization correction factor (icf) for an ion is the ratio of the total elemental

abundance divided by the ionic abundance. Because of collisional excitation and

recombination in H II regions, ionic abundances for our purposes are weighted by the

electron density. Thus,

!

X+i

X= icf "1 =

NX+iNedV#NXNedV#

$NX+iNedV#

NX NH( ) Np Ne( ) Ne2dV#

where NX/NH is the abundance of the element with respect to hydrogen by number, and

the simplifying assumption is made that Np/Ne is a constant, estimated from radio line

measurements. For two elements, X and Y, and/or two ionic states, +i and +j,

!

X+iX

Y+ jY

=NX+iNedV NX NH( ) Np Ne( ) Ne

2"[ ]"

NY+jNedV NY NH( ) Np Ne( ) Ne2"[ ]"

.

Thus, for lines at two wavelengths

!

" X+i( ) and

!

" Y + j( ) ,

!

X +i

Y + j"F# X+i( ) $# X+i( )

F# Y + j( ) $# Y + j( )

"NX+iNedV%NY+jNedV%

=X

+i/X

Y+ j/Y

&NX

NY

.

The derived ionic ratio,

!

X+iY

+ j is equal to the inverse ionization correction factors

!

X+iX Y

+ jY times the abundance ratio

!

NXNY

(Simpson et al. 1995).

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2b. Determine electron density, Ne and electron temperature, Te: Average electron

densities can be estimated from the ratio of the fluxes (correction for extinction) from two

lines of the same ion, which is equal to the inverse of the emissivity ratio for the two ions.

These ion pairs are density-sensitive due to the fact that the critical density for collisional

de-excitation is a strong function of wavelength, and the FIR lines suffer significant de-

excitation at the densities typically found in compact H II regions (Simpson et al. 1995).

We may estimate the electron density for our sources using the line pairs, [O III]

(52 µm/ 88 µm), [Ne III] (15.6 µm/ 36 µm) and [Ar III] (9.0 µm/ 21.8 µm) and an

existing model relating the chosen emissivity ratio (i.e. the inverse of the selected flux

ratios) to electron density, such as used by Lester et al. (1987). Figure 4 illustrates the

relationship between [O III] emissivity ratios, Ne and Te. Electron densities can also be

derived from existing radio continuum fluxes observed for these sources. A Te of 5000 K-

- assumed for Galactic center regions (Shaver et al. 1983) is assumed for this study.

Emissivities for each line can be calculated once Te and Ne are determined.

2c. Determine ionic abundances relative to hydrogen or relative to each other

(e.g. N++/O++). For optically thin FIR lines, ionic abundance relative to hydrogen (H or

H+) is given by the relation,

!

NXi

NH+

=F"

S#

3.485 $10%16T4

%0.35#5

%0.1

&"

Ne

N p

'

( ) )

*

+ , , (Rudolph et al. 1997)

where

!

NXi

and

!

NH+are the ion and proton abundances, respectively,

!

F" is the FIR line

flux in units of ergs sec-1cm-2,

!

S" is the free-free flux obtained from radio flux

measurements at frequency ν in units of Janskys, T4 is the electron temperature in units of

15

104 K, ν5 is the radio emission frequency in units of 5 GHz, and

!

"# is the emissivity per

unit volume of the FIR line at wavelength λ. The ratio of electrons to protons,

!

Ne N p , is

approximated by 1 +

!

NHe+

NH+( ) .

!

NHe+

NH+

will be approximated by using the plot of

!

He+H

+ vs. Teff from Rubin et al. (1988) (Rudolph et al. 1997).

2d. Determine the effective temperature (Teff) of the exciting star. We will use the

CLOUDY photoionization code to model the size of the H II region to determine the

ionization parameter, which is proportional to Teff (as described above). Strömgren theory

assumes that the number of recombinations equals the number of reionizations, and the

size of the H II region, is defined by,

!

rs"

3N

4#$

%

& '

(

) * 1/ 3

nH

+2 / 3 (Carroll and Ostlie 2007)

where rs is the Strömgren radius, α is the quantum-mechanical recombination coefficient

that describes the likelihood that an electron and a proton can form a hydrogen atom

(αnenH is the number of recombinations per unit volume per second) and N is the total

number of Lyman continuum photons produced by the star per second (Carroll and Ostlie

2007). We can then calculate Teff from the luminosity (

!

L = 4"Rstar

2

#Teff

), which in turn is

calculated through the relation of the Planck function and emissivity of the total Lyman

continuum.

2e. Correct the ionic abundance for unmeasured ionization states to obtain the

final abundances. The final step is to correct the ionic abundances relative to the

hydrogen abundance for the unmeasured ionization states, also called the icf, or

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ionization correction factor. The photoionization code, CLOUDY, will be used to

determine the ionization equilibrium temperature, Teff, by detailed modeling. The Saha

equation can then be used once Teff and Ne are obtained. The Saha equation relates atoms

in different states of ionization:

!

Xi+1

Xi

=2Zi+1

NeZi

2"mekT

h2

#

$ %

&

' ( 3/ 2

e)*ikT (Carroll and Ostlie 2007)

where

!

"iis the ionization energy needed to remove an electron from an atom or ion in the

ground state, thus taking it from ionization stage i to stage (i + 1), and Z is the partition

function (the weighted sum of the number of ways the atom can arrange its electrons with

the same energy):

!

Z = g je" E j "E1 kT( )

j=1

#

$ (Carroll and Ostlie 2007).

The appropriate partition functions can be determined using the modeled Teff. Thus, for

an observed fine-structure line, knowledge of all possible ionization states of the atom,

and other derived parameters, the Saha equation is used to calculate observed ionic

abundance relative to total abundance.

Once the total abundance is determined, the metallicity with respect to the

element in question can be calculated.

VII. Conclusions

We aim to use accurate abundance ratios for the inner 100 pc of the Galactic

center derived from this study, as well as a possible gradient in this inner region, to place

the Galactic center in the context of current understanding of metallicity across the

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Galaxy. An evaluation of metallicity indicators of enhancement processes in the Galactic

center via an analysis of N/O, (Ar,Ne)/(N,O), and (N,O,Ar,Ne)/H will also be explored.

VIII. References Afflerbach A, Chruchwell E. and Werner M. W. (1997) Galactic abundance gradients from infrared fine-structure lines in compact H II regions. The Astrophysical Journal 478, 190-205. Aller L. H. (1987) Physics of Thermal Gaseous Nebulae (Physical Processes in Gaseous Nebulae). Astrophysics Science Library, volume 112. D. Reidel Publishing Company. Carroll B. W. and Ostlie D. A. An Introduction to Modern Astrophysics, second edition. 2007, Addison-Wesley. Deharveng L., Peña M., Caplan J. and Costero R. (2000) Oxygen and helium abundances in Galactic H II regions—II. Abundance gradients. Monthly Notices of the Royal Astronomical Society 311, 329-345. Dwivedi B. N. and Gupta A. K. (1994) On the temperature measurement from the O IV emission lines. Solar Physics 155, 63-68. Garnett D. R. (1990) Nitrogen in irregular galaxies. The Astrophysical Journal 363, 142- 153. Giveon U., Sternberg A., Lutz D, Feuchtgruber H. and Pauldrach A. W. A. (2002) The excitation and metallicity of galactic H II regions from Infrared Space Observatory SWS observations of mid-infrared fine-structure lines. The Astrophysical Journal 566, 880-897. Güsten R., Walmsley C. M., Ungerechts H. and Chruchwell E. (1985) Temperature determinations in molecular clouds of the Galactic center. Astronomy and Astrophysics 142, 381-387. Güsten R. and Ungerechts H. (1985) Constraints on the sites of nitrogen nucleosynthesis from 15NH3-observations. Astronomy and Astrophysics 145, 241-250. Herter T. (1989). Infrared lines as probes of Galactic structure. Proceedings of the 22nd Eslab Symposium on Infrared Spectroscopy in Astronomy, Salamanca, Spain. ISO website (http://iso.esac.esa.int/) Lacy J. H., Achtermann J. M. and Bruce D. E. (1989) Observations of HI BR α, [Ne II] and [Ar III] from the central parsec of the Galaxy, in The Center of the Galaxy, (ed.) M. Morris, 523-524, IAU. Lester D. F., Bregman J. D., Witteborn F. C., Rank D. M. and Dinerstein H. L. (1981) The abundance of argon at the Galactic center. The Astrophysical Journal 248, 524-527. Lester D. F., Dinerstein H. L., Werner M. W., Watson D. M., Genzel R. and Storey J. W. V. (1987) Far-infrared measurements of N/O in H II regions: evidence for enhanced CN processing nucleosynthesis in the inner galaxy. The Astrophysical Journal 320, 573-585. Martín-Hernández N. L., Peeters E., Morisset C., Tielens A. G. G. M., Cox P., Roelfsema P. R., Baluteau J.-P., Schaerer D., Mathis J. S., Damour F., Chruchwell E. and

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Kessler M. F. (2002) ISO spectroscopy of compact H II regions in the Galaxy. Astronomy and Astrophysics 381, 606-627. Morris M. (1993) Massive star formation near the Galactic center and the fate of the stellar remnants. The Astrophyisical Journal 408, 496-506. Morris M. (2005) Massive star formation in the Galactic center. IAU Symposium 227, Acireale, Italy. Oort J. H. (1977) The Galactic center. Annual Review of Astronomy and Astrophysics 15, 295-362. Paumard T., Genzel R., Maillard J. P., Ott T., Morris M., Eisenhauer F., Abuter R. (2004). Census of the Galactic centre early-type stars using spectro-imagery. arXiv:astro-ph/0407189v1. Renzini A. and Voli M. (1981) Advanced evolutionary stages of intermediate-mass stars. Astronomy and Astrophysics 94, 175-193. Rubin R. H. (1985) Models of H II regions: heavy element opacity, variation of temperature. The Astrophysical Journal Suppl. Series. 57, 349-387. Rubin R. H., Simpson J. P., Erickson E. F. and Haas M. R. (1988) Determination of N/O from far-infrared line observations of Galactic H II regions. The Astrophysical Journal 327, 377-388. Rubin R. H., Simpson J. P., Haas M. R. and Erickson E. F. (1991) Axisymmetric model of the ionized gas in the Orion nebula. The Astrophysical Journal 374, 564-579. Rudolph A. L., Simpson J. P., Haas M. R., Erickson E. F. and Fich M. (1997) Far- infrared abundance measurements in the outer Galaxy. The Astrophysical Journal 489, 94-101. Shaver P. A., McGee R. X., Newton L. M., Danks A. C. and Pottasch S. R. (1983) The Galactic abundance gradient. Monthly Notices of the Royal Astronomical Society 204, 53-112. Simpson J. P., Colgan W. J., Rubin R. H., Erickson E. F. and Haas M. R. (1995) Far- infrared lines from H II regions: abundance variations in the Galaxy. The Astrophysical Journal 444, 721-738. Thuan T. X., Izotov Y. I. and Libovetsky V. A. (1995) Heavy element abundances in a new sample of low-metallicity blue compact Galaxies. The Astrophysical Journal 445, 108-123. Wannier P. G. (1989) Abundances in the Galactic center in The Center of the Galaxy, ed. M. Morris, 107-119, IAU. Willner S. P., Russell R. W., Puetter R. C., Soifer B. T. and Harvey P. M. (1979). The 4 to 8 micron spectrum of the Galactic center. The Astrophysical Journal 229, L65- L68. Wink J. E., Wilson T. L. and Bieging J. H. (1983) An H76α survey of Galactic H II regions: electron temperature and element gradients. Astronomy and Astrophysics 127, 211-219.

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IX. Figures and Tables

Figure 1. 2MASS image of the Arches massive- star forming region ~ 30 pc from the galactic center (slide image from Morris 2005).

Figure 2. Whirlpool galaxy (M 51) showing H II regions (red patches). Red color is from H recombination emission. Companion galaxy NGC 5195 is also shown (http://sci.esa.int/sciencee/www/object/index.cfm?fobjectid=37004).

20

Table 1. Accessible fine-structure lines probed by the Short-Wave Spectrometer. Fluxes for H II metallicity region G-0.02.0.07 are shown (this study).

Table 2. Accessible fine-structure lines probed by the Long-Wave Spectrometer. Fluxes for H II metallicity region G-0.02.0.07 are shown (this study).

21

Figure 3. Short-wavelength emission lines for metallicity region G-0.02-0.07 (this study).

Figure 4. Volume emissivity ratio of [O III] 51.8 µm as a function of electron temperature and electron density (Lester et al. 1987).