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Observations of interstellar H2: Milestones

It was suggested for quite a long time that hydrogen in its molecular form might well exist in some parts of the interstellar space, most likely in dense interstellar clouds which shield their interior from energetic, H2 dissociating radiation (e.g., Eddington 1937; Gould & Salpeter 1963; Gould et al. 1963; Hollenbach et al. 1971; but see Strömgren 1939). Up to the late 1960s however, H2 was only observed in the atmospheres of planets and some stars (see reviews by Field et al. 1966; Shull & Beckwith 1982). The first signatures of interstellar H2 were found through rocket and satellite observations as UV absorption features (due to electronic transitions) in the spectra of some bright stars (Carruthers 1970; Smith 1973; Spitzer et al. 1973), and had their origin in thin interstellar clouds in the line of sight to the stars. The detection of electronic transitions at UV wavelengths and of pure rotational transitions at mid-infrared wavelengths both require air-based or space-borne observatories. In contrast, many ro-vibrational transitions occur in atmospheric windows at far-red and near-infrared wavelengths and are thus observable with comparatively low technical efforts. A search for those lines was suggested as an interesting alternative (Gould & Harwit 1963; see also Osterbrock 1962). However, the first attempts to detect these lines at far-red optical wavelengths in dark clouds near hot stars proved difficult (Werner & Harwit 1968; Gull & Harwit 1971). Finally, a number of ro-vibrational H2 emission lines were found in K-band spectra of the planetary nebula NGC 7027 (Treffers et al. 1976) and towards the BN and KL regions in Orion (Gautier et al. 1976; Grasdalen & Joyce 1976; Beckwith et al. 1978a; 1979). It became clear very soon that the observed lines in Orion (characteristic of a ~2000 K warm gas, apparently coming from only a rather small mass of H2 gas) could not represent the bulk of the H2 gas in the region. Instead, heating of the gas in a shock wave was proposed (e.g., Hollenbach & Shull 1977; Kwan 1977; London et al. 1977; Draine & Roberge 1982; Chernoff et al. 1982).

The following years brought a number of other detections of ro-vibrational H2 sources, many of them in star forming dark clouds (see compilation in Shull & Beckwith 1982). These included H2 emission from (or from around) the prototype T Tauri star T Tau (Beckwith et al. 1978b), a number of intermediate to high mass star forming regions (DR 21, OMC-2: Fischer et al. 1980a; NGC 7538: Fischer et al. 1980b; NGC 2071, Cep A, GL 961: Bally & Lane 1982; W51: Beckwith & Zuckerman 1982; NGC 6334: Fischer et al. 1982), and a number of Herbig-Haro objects (HH 1 & 2, HH 46, HH 53, HH 54: Elias 1980; see also Fischer et al. 1980a). The generally observed coexistence of H2 emission regions with high velocity molecular CO outflows (e.g., Fischer et al. 1985; Bally & Lane 1982; Simon & Joyce 1983; Burton et al. 1989b), the spectral properties, and the detection of H2 emission in Herbig-Haro objects suggested that the H2 emitting regions trace shock heated gas in outflows from young stellar objects, similar or equivalent to the optically visible Herbig-Haro objects.

This suggestion was supported by further observations (e.g., Taylor et al. 1984; Zealey et al. 1984, 1986; Lane & Bally 1986; Garden et al. 1986; Sandell et al. 1987; Schwartz et al. 1987; Zinnecker et al. 1989; Wilking et al. 1990a) with increasing sensitivity and spatial as well as spectral resolution (particularly after the installation of infrared array cameras; e.g., Schwartz et al. 1988; Hartigan et al. 1989; Lane 1989; Garden et al. 1990). However, it also became clear that the H2 shocks in some cases had to originate in different parts of the shock fronts than optical Herbig-Haro objects. These sometimes have shock velocities of order 200 km/s (e.g., HH 1/2, Hartigan et al. 1987), whereas H2 molecules should only survive in shocks with a velocity up to ~50 km/s (see below). The solution to this problem is that H2 molecules can survive in parts of fast moving shock waves, where the shock front is not parallel to the direction of propagation of the shock wave, e.g., in the wings of bow shocks (see Fig. 7). There the velocity component of the shock front along the direction of propagation is much smaller. Besides emission from bow shock like working surfaces (whether internal or at the leading working surface), in some cases turbulent mixing or shear layers along the jet beam or outflow cavity walls may be responsible for H2 emission (e.g., HH 26A, HH 40: Chrysostomou et al. 2000; Davis et al. 2000a; Zinnecker et al. 1989; HH 46/47: Eislöffel et al. 1994a; see also Noriega-Crespo 1997). In a few cases, H2 emission may also originate in ``shocked cloudlets'' immersed in the flow (e.g., HH 11: Davis et al. 2000a, see also Hartigan et al. 1987).

In the following, only a few selected observations will be highlighted, but a much larger body of examples of H2 emission in flows from young stars exists (see Eislöffel 1997 for an overview of recent H2 observations in flows from young stars).


next up previous contents
Next:Origin of the H2 emission: excitation mechanisms Up:2.3 Molecular hydrogen jets Previous:The H2 molecule

Thomas Stanke

Mon Aug 14 17:51:08 CEST 2000
Last modified Thu Jul 12 15:24 CEST 2001