Stephen S Eikenberry

7.1. Introduction

Ultimately, the overwhelming majority of emission-line sources in the Universe are "galactic sources" - meaning discrete objects located within a particular galaxy (rather than some global property of a galaxy or some source not located in a galaxy). However, the most common of these, H ii regions, are so ubiquitous that they are being covered elsewhere in this volume as the "baseline" source of emission lines. In addition, most of the other chapters are devoted to line emission either integrated over entire galaxies (or significant portions thereof) or from active galactic nuclei.

Given that coverage, I will focus this chapter primarily on "stellar" sources of line emission in the Milky Way other than HII regions - including young stellar objects, massive and/or evolved stars, and stellar remnants (planetary nebulae, supernova remnants, and accreting compact objects in binary systems). I will also put considerable emphasis on emission lines with rest wavelengths in the near-infrared waveband, due to the importance of this waveband for probing the dusty planar regions of the Milky Way where most of these sources are to be found.

In the sections below, I will begin with a review of important diagnostic optical emission lines and a more-detailed overview of key (rest-wavelength) infrared emission lines. I will then move on to "nebular" sources of emission lines (omitting HII regions, but including planetary nebulae and supernova remnants). Next, I will cover "stellar" sources (including young stellar objects, "main-sequence" massive stars, and evolved Wolf-Rayet stars). After that, I will move on to emission lines arising from accreting stellar remnants (cataclysmic variables, neutron-star X-ray binaries, and black-hole X-ray binaries).

7.2. Overview of infrared emission lines

As noted above, other chapters in this volume address some of the primary diagnostic optical emission lines associated with HII regions. These include the Balmer series of hydrogen emission lines, as well as OII, NII, and SII. However, rest-wavelength infrared (IR) lines, particularly those in the wavelength range of 1.0 — 2.5 |im, while less widely known to most of the astronomical community, are of critical importance for studying Galactic sources of line emission. In the following subsections, I will review why IR emission is important for Galactic objects, and observational aspects for IR line emission (detectors/instruments, IR atmospheric windows, IR atmospheric emission). I will then move on to review key emission-line features from hydrogen (particularly Paschen and Brackett series), heavier-element emission, and particularly molecular line emission.

7.2.1 Why infrared emission lines?

Many readers may ask the question, "Infrared - why bother?" (In fact, many professional astronomers feel that way even now.) From a simple atomic-physics perspective, the IR looks less than promising. Hydrogen, the most-abundant element in the Universe, has its strongest transitions - the Lyman series - in the ultraviolet (UV) waveband, rather far from the IR. Of course, then, one could ask "Why the optical?," given that the Balmer

The Emission-Line Universe, ed. J. Cepa. Published by Cambridge University Press. © Cambridge University Press 2009.

series is intrinsically weaker than Lyman lines. The primary answer to that question is that the Universe in general (and the Earth's atmosphere in particular) is not transparent to Lyman radiation. Since most cosmic hydrogen is in its ground state, it will absorb Lyman photons. Interstellar dust also wreaks havoc on these short wavelengths, both for Lyman series and for other non-hydrogen transitions. As a result, optical lines such as the Balmer series are often more powerful probes than the Lyman series (at least in the local Universe) despite their weaker atomic strengths.

While the Earth's atmosphere is less "friendly" to the IR waveband than it is to the optical (see Section 7.2.2), "cosmic transparency" is even better for the IR - to the point that for many Galactic situations it is the only waveband of the UV, optical, and IR set which is useful for observing emission-line sources. Again, the primary reason for this is the simple physics of radiative processes. With typical grain sizes of <1.0 |m, interstellar dust grains simply do not have much cross-section for scattering or absorbing photons at wavelengths A > 1.0 |im. These dust grains are of course strongly concentrated in the plane of the Milky Way, as are most star-forming regions, massive stars, planetary nebulae, supernovae remnants, etc., making the IR waveband invaluable for studies of these objects.

Figure 7.1 shows a view of our Galactic Center as seen in the "optical" B, R, and I bands, as compared with the near-IR J (1.1 — 1.4 |im), H (1.5—1.8 |im) and Ks (1.95—2.4 |im) bands. While the optical pictures are actually deeper than the IR in terms of magnitudes, the vastly superior penetrating power of the IR (AK ~ 0.1 AV - a factor of ten in magnitudes!!) reveals many more sources of emission. For the Galactic Center, AK ~ 3 mag, meaning that >6% of the emitted IR photons make it through the dust, while AV ~ 30 mag, meaning that only one optical photon of 1012 emitted makes it through.

7.2.2 Observing in the infrared waveband While we have just seen that "longer is better" for observing wavelengths, at least in terms of Galactic dust extinction, very important non-cosmic effects usually limit sensitive IR observations to A < 2.5 |m (from both space and the ground, for now). There are several factors that contribute to this approximate limit, which I will now review.

7.2.2.1 Infrared observing technology

In recent years, IR detector and instrument technology has seen tremendous gains, and now closely resembles optical detection technology in many ways. However, there are important - even critical - differences between the two, which tend to make IR instrumentation much more challenging than optical. These differences all spring from the fundamental fact that in this waveband - 1.0—2.5 |m - the outstanding optical detection technology of the silicon-based charge-coupled device (CCD) is no longer functional. The energies of individual IR photons are simply too low to excite valence-band electrons across the silicon bandgap into the conduction band where they can be collected and "detected," or measured. This requires erstwhile IR observers to seek out and use alternative semiconductor devices for their instrumentation.

The IR detectors of choice for this waveband currently come in two primary flavors -HgCdTe and InSb arrays. Owing to the trillions of dollars invested in silicon technology over the past few decades, neither of these can match the cosmetic quality, large format, and relatively low cost of CCD detectors. HgCdTe is a ternary alloy and its tunable composition allows one to vary its bandgap energy and thus cutoff wavelength for detection. The most well-developed HgCdTe mixture, in terms of cosmetic quality, array format, quantum efficiency, noise properties, and simple number of IR arrays in

Hgcdte Infrared Detectors
Figure 7.1. Multi-wavelength views of the Galactic Center region in the DPOSS B band (top left); DPOSS R band (top right); DPOSS IR band (middle left); 2MASS J band (middle right); 2MASS H band (bottom left); and 2MASS Ks band (bottom right).

use, has a cutoff wavelength of 2.5 |m. The current state-of-the-art HgCdTe arrays have 2048 x 2048-pixel formats, quantum efficiency of ~ 70%, and reading noise of about ten electrons. InSb technology also produces high-quality arrays, with a bandpass cutoff of ~ 5.5 |im. InSb arrays have similar formats to HgCdTe arrays, and higher quantum efficiency (>90% - nearly perfect). However, the advantage of the latter can be offset by the higher reading noise (about 25 electrons) and the disadvantage (for ground-based applications) of the longer cutoff wavelength.

In addition to the relatively poorer quality of IR detectors compared with CCD technology, their longer-wavelength cutoff also enforces a crucial global design requirement -cryogenic cooling. If the detectors have temperatures such that kT ~ Ebandgap, then the thermal distribution of electron energies will excite electrons over the semiconductor bandgap to the conduction band, making them appear identical to photoelectrons - a primary source of "dark current" in detectors. Since Ebandgap = he/Xc, where Ac is the cutoff wavelength of the detector material, to avoid overwhelming dark current the detectors must be cooled to T ^ hc/kAc. For HgCdTe with a 2.5-|im cutoff, this corresponds to T ~ 65 K, whereas for InSb with a 5.5-|im cutoff, it corresponds to T ~ 30 K. Similarly, any optical materials (lenses/mirrors) and their support structures must also be cooled enough to prevent the Wien tail of their blackbody glow from swamping the detectors (albeit, this requirement usually allows slightly higher temperatures than for the detectors themselves). These temperatures require large and cumbersome vacuum/cryogenic systems for IR instruments, in addition to unusual, difficult, and expensive optical and structural materials.

7.2.2.2 The infrared atmosphere

In addition to the differences in instrument-technology requirements, another major issue for IR observations (at least from the ground) is the Earth's atmosphere. While the Earth's atmosphere is largely transparent and dark in the optical waveband, it is both partially opaque AND largely bright in the IR (!). Figure 7.2 shows the atmospheric transmission curve for the Earth over the UV-IR bandpass.

In the waveband range 1.0-2.5 |im, the atmosphere transmits well only in certain atmospheric "windows." These constitute the near-IR bandpasses of the J(1.1 —1.4 |m), H(1.5—1.8 |m), and K(1.95—2.45 |m) bands. Note that there are also windows at longer wavelengths than this, which are even less susceptible to dust extinction than JHK bands. However, at these wavelengths, thermal emission at temperatures commonly found in the Earth's atmosphere dominates the cosmic "sky background," making sensitive observations at these wavelengths very difficult from the ground (with the possible exception of from very cold locations such as the South Pole).

While thermal atmospheric emission is typically small in the JHK bandpass (except on warm nights at the long end of the K band - which has led to the common use of the shorter-cutoff Ks filter), the atmosphere is still much brighter than in the optical waveband. In the J and H bands, as well as the short part of the K band, this is primarily due to a veritable forest of telluric OH emission lines (see Figure 7.3), while thermal emission begins to become non-negligible in the long-wavelength end of the K band.

7.2.3 Important infrared emission lines

While Galactic objects certainly have optical emission, and much of the observational work to date revolves around them, those lines will be presented and reviewed elsewhere in this volume. There are several important sets of IR emission lines that are key for studies of Galactic sources that have become extinct, which will not be reviewed elsewhere, so I will go over them here.

Atmospheric Transmission

Figure 7.2. An atmospheric-transmission curve for Mauna Kea (adapted from Gemini

Observatory website).

Figure 7.2. An atmospheric-transmission curve for Mauna Kea (adapted from Gemini

Observatory website).

7.2.3.1 Hydrogen emission in the infrared

The primary "infrared" hydrogen emission series are the Paschen and Brackett series, with their lowest energy levels being 3 and 4, respectively. This is to be compared with the "ultraviolet" Lyman series (lowest energy level 1) and the "optical" Balmer series (lowest energy level 2). Table 7.1 summarizes the key transitions for each.

An important factor to note from the above is that none of the strongest "a" transitions for these series lies in a wavelength range easily observable from the ground. The strongest Paschen-series line observable from the ground is the Paß transition, but for Brackett series the strongest is Bry. Similarly, all but the transitions near the continuum limit

Table 7.1. Hydrogen emission lines in the infrared

Line

Transition

Wavelength (^m)

Paa

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