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Unlocking the Keyhole - H2 and PAH emission from molecular clumps in the Keyhole Nebula

To better understand the environment surrounding CO emission clumps in the Keyhole Nebula, we have made images of the region in H2 1-0 S(1) (2.122 um) emission and polycyclic aromatic hydrocarbon (PAH) emission at 3.29 um. Our results show that the H2 and PAH emission regions are morphologically similar, existing as several clumps, all of which correspond to CO emission clumps and dark optical features. The emission confirms the existence of photodissociation regions (PDRs) on the surface of the clumps. By comparing the velocity range of the CO emission with the optical appearance of the H2 and PAH emission, we present a model of the Keyhole Nebula in which the most negative velocity clumps are in front of the ionization region, the clumps at intermediate velocities are in it, and those which have the least negative velocities are at the far side. It may be that these clumps, which appear to have been swept up from molecular gas by the stellar winds from eta Car, are now being over-run by the ionization region and forming PDRs on their surfaces. These clumps comprise the last remnants of the ambient molecular cloud around eta Car.


INTRODUCTION
The Keyhole Nebula is part of the Carina Nebula, a star forming/giant molecular cloud (GMC) complex, containing some of the most massive stars known in our galaxy. It is bathed in the UV radiation field from the nearby open star cluster, Tr 16. This cluster contains numerous O-type stars, including three O3 stars (Walborn 1995). It is at a distance of 2.2 kpc (Tovmassian 1995) and is centred on one of the most massive stars known -η Car, which has a present day mass of about 120 M⊙ . The optical image of the Keyhole Nebula taken by Malin (1993) reveals many interesting features such as dark patches and filaments, the most prominent being in the shape of an oldfashioned key hole, which gave the nebula its name (see Fig. 1a). Herschel (1847) was the first to note this peculiar feature and likened it to the shape of a "lemniscata". Some of the smaller features also have bright rims. This is beautifully illustrated in the recent colour-composite image taken with NASA's Hubble Space Telescope. † The molecular gas in this region is highly inhomogeneous, breaking into several clumps which are all associated with the dark optical features . The clumps, typically a few tenths of a pc across (10-30 arcsec) and with masses of ∼ 10 M⊙ are distributed across a range of velocities (see Table 1). Some of the molecular clumps correspond to those optical features that have bright rims, suggesting the presence of photodissociation regions (PDRs) at the interface with the surrounding ionized gas. PDRs are regions where the far-UV radiation field dominates the heating and chemistry of the gas. They begin at the sharp ionized/neutral interface at the edge of molecular clouds and then extend well into the molecular gas (see Tielens et al. 1993).
To better understand the environment surrounding these clumps we have imaged the Keyhole region at two wavelengths suitable for detecting emission from PDRs.
(a) Three-colour optical emission image, courtesy of David Malin. This image was produced by combining 3 exposures (of 15 mins) each taken using blue, green and red filters. The Keyhole is the prominent dark band bisecting the image.
(b) H 2 2.122 µm line emission obtained with UNSWIRF. This image is continuum subtracted (aside from the circular residuals from η Car).
(c) PAH 3.29 µm emission obtained with SPIREX. This image includes continuum sources (the diffraction spikes from η Car are artifacts).  Tables 1 and 2.  Table 1. A list of the 12 CO(2-1) emission clumps measured by . The mass measurements may underestimate the true mass of each clump by up to a factor of 3. The type of dark optical feature each clump is associated with is also listed. The velocities are given with respect to the local standard of rest (LSR).

Clump Velocity
Mass We have used the University of New South Wales Infrared Fabry-Perot (UNSWIRF) to image H2 1-0 S(1) (2.122 µm) emission and the SPIREX/Abu thermal infrared camera at the South Pole to image narrow-band emission at 3.29 µm. This wavelength includes one of the more prominent 'unidentified' infrared (UIR) emission bands and is generally attributed to the C-H stretching mode of free polycyclic aromatic hydrocarbon molecules (PAHs) (Léger & Puget 1984, Allamandola & Tielens 1985, Geballe et al. 1994, Joblin et al. 1995. H2 1-0 S(1) emission in molecular clouds can be attributed to collisional excitation in the presence of shocks or to UV fluorescence in PDRs (see Burton 1992). In dense gas, n ≥ ncrit ∼ 10 5 cm −3 where n is the hydrogen nucleus density, UV excitation levels are collisionally redistributed leading to an increase in the population of the ν = 1 level. This subsequently increases the intensity of the ν=1-0 S(1) line (see Allen et al. 1999). Previous observations of PDRs indicate that 3.29 µm PAH emission zones are situated near the very edge of molecular clouds, peaking at the ionization front, while H2 1-0 S(1) emission zones extend further into the clouds by about 2 magnitudes of visible extinction (e.g. Sloan et al. 1997, Tielens et al. 1993, Burton et al. 1989. So while H2 1-0 S(1) emission cannot be used alone to identify PDRs, if found to correspond to another PDR tracer (such as PAH emission) it can pinpoint regions of molecular gas that are high in density and exposed to the surrounding far-UV radiation field.

UNSWIRF observations
UNSWIRF (Ryder et al. 1998) is a near-infrared tunable imaging spectrometer. It was used in conjunction with the Infrared Imager and spectrometer (IRIS) on the Anglo-Australian Telescope (AAT), and produces a circular field of 1.7 arcmin diameter at 0.77 arcsec pixel −1 .
Images of H2 1-0 S(1) (2.122 µm) emission towards the Keyhole region were obtained on 1996 April 24. A series of 30 pointings were made on a 1 arcmin grid at the wavelength of the 2.122 µm line. An integration time of 60 s was used. An image of the sky was obtained after every five minutes, offset by 5 arcmin from the previous position. This entire sequence was repeated at 3-4 profile widths from the line centre in order to sample the corresponding continuum emission. Bias subtraction and linearisation were performed during readout. All the object images were sky subtracted and then flat-fielded using normalised dome flat-fields. The continuum images, appropriately scaled, were then subtracted to leave just pure line emission in each. However, the strong emission from η Car prevented a perfect subtraction and so residuals appear in the final image in frames containing η Car. The continuum images were then combined into a mosaic sequence. A coordinate axis was added using a corresponding image obtained with the Digitised Sky Survey (see see http://skyview.gsfc.nasa.gov/cgi-bin/surv comp.pl?dss) and the program koords (Gooch 1996). The same mosaic sequence and coordinate axis were then applied to the emission-line images. Observations of the spectroscopic standards BS 4013 and HD 105116 were used for the flux calibration and continuum scaling factors. The central wavelength measured varies by up to ∼ 10 km s −1 between pixels (∼ 1/7 of the FWHM). This can be corrected for by the use of a calibration lamp at closely-spaced etalon settings. Normally line images are then obtained by measuring at several plate spacings (typically 3-5) and fitting the instrumental profile to the line. However, because of the considerable overheads this incurs, this was not attempted for the mapping mode of operation used here. We estimate this introduces an uncertainty of ∼ 30 per cent in the line fluxes. The seeing throughout the observations was 1.5 arcsec which, combined with tracking errors and co-addition of frames, produced a positional uncertainty of ± 2.5 arcsec. The final image has a 1σ rms of 1 × 10 −16 erg s −1 cm −2 pixel −1 .

SPIREX/Abu observations
SPIREX is a 60-cm telescope commissioned in 1994 for operation at the South Pole (Herald et al. 1990). It was equipped in 1998 with the NOAO Abu IR camera (Fowler et al. 1998), which incorporates an engineering grade 1024 × 1024 'Aladdin' InSb detector array able to image a circular field of view of diameter 10.2 arcmin with a 0.6 arcsec pixel scale.
Observations were carried out on 1998 November 11 (during daylight) using the narrow-band PAH filter. The response curve of this filter covers 3.262 -3.336 µm and is centred on 3.299 µm, with a half-power width of 0.074 µm. An integration time of 30 s was used for each frame. The overlapping region from a total of 52 frames was used to obtain a final image of 9 arcmin in diameter. The observing sequence consisted of 5 sky frames followed by 10 object frames. All frames were dark subtracted and flat-fielded. For the sky subtraction, a scaled median of the 6 nearest frames, regardless of their actual image type, were subtracted from each object image. The final image contains both PAH emission and continuum emission. A coordinate axis was added using a corresponding image obtained with the Digitised Sky Survey and the program koords.
Observations of the standard star, δ Dor, observed 1998 November 19, were used for flux calibration (L band magni-tude of 3.711 mag). We estimate the PAH fluxes to have an uncertainty of ±5 per cent. The diffraction limit at 3.3 µm of the telescope is 1.4 arcsec, comparable to the typical icelevel seeing of the site. However, a combination of tracking errors, tower shake and co-addition of frames limited the pointing accuracy to ± 2.8 arcsec. The final image has a 1σ rms noise of 2 × 10 −15 erg s −1 cm −2 pixel −1 .
The observed wavelength range includes the Pf δ hydrogen recombination line. From the peak Br γ emission flux measured towards the southern edge of clump A1 (6 × 10 −15 erg s −1 cm −2 arcsec −2 ) and the calculated relative intensities of the Pf δ and Br γ emission lines (0.26 for the case of Ne = 10 4 cm −3 and Te = 10 4 K, as listed in Table 6 of Hummer & Storey 1987), we estimate the emission from Pf δ is less than 2 × 10 −15 erg s −1 cm −2 arcsec −2 . This value is the same as the 1σ rms uncertainty of the final 3.29 µm image and therefore we assume that the contribution from the Pf δ line is negligible. Fig. 1(b) shows a map of the H2 emission towards the Keyhole Nebula. The emission is distributed into seven discrete clumps of less than 1 arcmin in size spread over an area of 4 × 4 arcmin 2 . Each emission feature corresponds to a 12 CO(2-1) emission clump identified by  and has been labelled accordingly, with the exception of the faint H2 emission clump located in the north-eastern edge of the image which we have labelled U1. No CO observations have been made towards this clump. Fig. 1(c) shows an image of the PAH emission over the same region of sky. The emission is also distributed into clumps, all of which (except clump U1) can be matched almost exactly with H2 emission clumps. This confirms the H2 emission is caused by UV fluorescence and that the clumps are surrounded by PDRs. Table 2 lists the peak position, peak intensity, total integrated intensity and size of each H2 emission clump as well as the corresponding PAH peak intensity and total integrated intensity. Fig. 2 shows the PAH emission in grey-scale overlaid with contours of H2 emission for clumps A1 ‡ , A2, 4, 5, B and C . The overall distribution of the H2 and PAH emission is similar for all the clumps, however, there are smallscale differences. The emission is not uniform but exhibits many small bright knots and in some cases the H2 emission knots do not correspond to PAH emission knots. This is to be expected considering that the H2 and PAH emission zones within a PDR may not completely coincide with one another. The bright H2 and PAH emission knots are pinpointing high-density regions at different UV optical depths within the molecular clumps.

RESULTS
The ratio of the PAH total intensity to the H2 1-0 S(1) total intensity is reasonably consistent from clump to clump. Clump A2 has the highest ratio (44), while clump B has the lowest (20). ‡ We have dubbed clump A1 the 'Kangaroo Nebula' owing to its morphological resemblance to an Antipodean marsupial, see Ryder et al. (1998). When comparing the H2 and PAH emission with the optical emission of Fig. 1(a) it appears that all the emission clumps correspond to dark optical features. Clumps A1, A2, 4 and 5 correspond to prominent patches which have bright rims, while clumps B and C coincide with very faint, filamentary patches. It is worth noting that no H2 or PAH emission was detected towards the most prominent optical obscuration-the Keyhole itself, or towards the boomerangshaped filament to the western edge of the image. Neither of these latter two dark features appears to have a bright rim, unlike clumps A1, A2, 4 and 5.
The Keyhole has been traced in 12 CO(2-1) emission by  and corresponds to clumps 1, 2, and 3 (see Table 1). These three clumps are the only molecular clumps detected by  which show no H2 or PAH emission.
As listed in Table 1, the velocities with respect to the local stand of rest (LSR) of the molecular clumps taken from their 12 CO(2-1) emission vary from −32 to −8 km s −1 . When comparing these velocities with the clump's corresponding optical, PAH and H2 emission features there appears to be an interesting trend. Firstly, the clumps with the most negative velocities (clumps 1, 2, and 3) are those which have the most prominent optical obscuration and also are the only clumps with no associated PAH and H2 emission. Secondly, those clumps with intermediate velocities (clumps  a H 2 image has a 1 σ rms of 0.1 × 10 −15 erg s −1 cm −2 arcsec −2 b PAH image has a 1σ rms of 3 × 10 −15 erg s −1 cm −2 arcsec −2 c 3σ upperlimit A1, A2, 4 and 5) correspond to the optical features with bright rims. Thirdly, the clumps which have the least negative velocities (clumps B and C) correspond to the very faint patches. Using these results we present a model for the Keyhole Nebula whereby the velocities of the clumps are associated with positions along our line of sight; clumps 1, 2, and 3 (the Keyhole) are in front of the nebula, clumps A1, A2, 4 and 5 are in it and clumps B and C are located at the far side. This model is represented schematically in Fig. 4. Radio recombination-line observations of the ionized gas in this region (Car II) reveal a complex velocity structure which has been interpreted as an expanding shell centred close to η Car (Huchtmeier & Day 1975, Deharveng & Maucherat 1975. There are two main velocity components whose relative fluxes vary across the region; one having LSR velocities between −6 to −17 km s −1 ; and the other having LSR velocities between −23 to −41 km s −1 . These velocities are consistent with clumps A1, A2, 4 and 5 (i.e. those with bright rims), which have LSR velocities between −17 and −25 km s −1 , being located in the ionized region and the remaining clumps, with LSR velocities near −8 and −30 km s −1 , being situated at the far and near side, respectively. The systematic LSR velocity of η Car is taken to be −18.5 ± 10 km s −1 , similar to the other O-type stars in the region , placing it in the middle of the range.

Dynamics
It is generally thought that the strong stellar winds from the massive stars belonging to Tr 16, in particular η Car, are responsible for the complex dynamics of the ionized and molecular gas in the Keyhole region (e.g. , Cox 1995and Deharveng & Maucherat 1975. Here we extend the analysis of previous studies to include a better estimate of the mechanical luminosity associated with η Car and the energy conversion efficiency factor required to sustain the kinetic energy of a typical clump. The kinetic energy of a typical clump, KC, with a mass of 10 M⊙ and a velocity of 10 km s −1 , is of the order 10 46 erg. We can equate this value with the mechanical luminosity, LW, of the stellar wind of η Car using, KC ∼ LW Ω τ ǫ, in order to obtain the required conversion efficiency factor, ǫ. We estimate the solid angle, Ω, of a typical clump of crosssection 1 arcmin located 3.5 arcmin from η Car is ∼ 0.06. The mechanical luminosity of the wind for η Car, using a mass-loss rate of 5×10 −4 M⊙ yr −1 ) and a wind velocity of 700 km s −1 (Lamers 1989) is ∼ 8 × 10 37 erg. If we take the time period, τ , to be equal to the passage time across the ionized nebula, ∼ 10 5 years (2 arcmin at 10 km s −1 ), then the efficiency factor need only be ∼ 0.01 per cent.
In practice τ is likely to be very much less than 10 5 years, as the clumps will rapidly decelerate as they sweep up material. We estimate its value following Dyson & Williams (1997). They show that the radius of a stellar wind-blown bubble is given by ∼ (LW/nτ 3/5 . For a typical clump-mass of 10 M⊙, the required density (n) of the swept-up gas is ∼ 1 000 cm −3 . This yields τ ∼ 10 3 years for the η Car wind which leads to a required efficiency factor of ∼ 1 per cent. This efficiency factor is still much less than the 20 per cent efficiency factor estimated by Dyson & Williams (1997) for this process for a steady wind.
Thus, the kinetic energy of the clumps can readily be explained by the mechanical energy of a wind from η Car. We conclude that the clumps represent swept-up ambient cloud material which is being driven out of the ionized region by such a wind. In this scenario we would expect the sweptup mass of each clump to be much greater than its wind mass (given byṀ Ω τ 4 π ). Substituting the figures above, we obtain 0.003 M⊙ for the wind mass, considerably less than the typical swept-up mass of 10 M⊙, as expected.

Excitation
The morphology of the H2 and PAH emission strongly suggests that UV fluorescence is responsible for the excitation. This conclusion can be quantitatively tested by comparing the observed fluxes with those expected for the physical conditions within the nebula.
We first estimate the ionizing (hv = 13.6 eV → ∞) and far-UV (hv = 6 eV → 13.6 eV) fluxes incident on clump A1 (taken to be 1 arcmin across), from both η Car (3.5 arcmin away) and the O3 star HD 303308 (2.8 arcmin away). HD 303308 is taken to have T eff ∼ 47 500 K and radius 1 × 10 12 cm, following Panagia & Walmsley (1978). Using Kurucz's (1997) model stellar spectra we find FUV = 0.34 erg s −1 cm −2 and FFUV = 0.26 erg s −1 cm −2 . The spectral type of η Car is uncertain. Based on the observations of  at wavelengths longer than 1 µm, its luminosity has been estimated by  to be 5×10 6 L⊙. At such high luminosities, blackbody and Kurucz models indicate that we can assume the bulk of the flux is distributed over the ionizing and far-UV wavelengths and that FUV ∼ FFUV. Using a stellar radius of 6 × 10 12 cm (Davidson & Humphreys 1997) we find that for η Car FUV ∼ FFUV ∼ 13 erg s −1 cm −2 , which is considerably greater than the corresponding fluxes from HD 303308.
We thus estimate the total number of ionizing UV pho-tons incident on clump A1 is NLycont ∼ 3 × 10 48 photons s −1 and are mostly from η Car. This number is an upper limit, and may be reduced if either internal extinction is present or the distance from the source is greater than its projected value. We have obtained an additional observation of Br γ emission at 2.166 µm using UNSWIRF. The total Br γ flux from a region covering 0.7 arcmin 2 , adjacent to the northwest-southeast edge of clump A1, is 3.9 × 10 −11 erg s −1 cm −2 . From this, assuming that each Lyman-continuum photon produces 70 Br γ photons, we derive NLycont ∼ 2 × 10 48 photons s −1 , remarkably similar to our estimate based on the stellar luminosities.
A typical clump with a surface area of 1 arcmin 2 , hydrogen molecule density of 10 4 cm −3 and mass of 10 M⊙ yields an emitting thickness, Av ∼ 1 (in a direction away from the exciting star) of ∼ 10 17 cm (4 arcsec at 2.2 kpc). This order of magnitude calculation implies widths similar to those observed for the elongated clumps, 4 and 5, suggesting that they are seen edge-on. These two clumps and the larger clumps, A1 and A2, are thought to be situated inside the H II region (see Fig. 4). Clumps 4 and 5 may be moving edge on, in the plane of the sky, and clumps A1 and A2 may be moving mostly face on, away from us.
The emission measure, n 2 e L, derived from the Br γ line flux is 6 × 10 5 cm −6 pc. Taking a typical clump thickness as representative of L gives L ∼ 4 arcsec (or 0.03 parsec based on the above estimate). This yields an electron density, ne ∼ 4 × 10 3 cm −3 for the ionized gas, which is intermediate in value between the estimates for the initial density and that swept-up in the clumps. The emission measure is also consistent with the value of 2 × 10 5 cm −6 pc obtained from measurements of the 4.8 GHz continuum emission over the same region (Brooks et al., in preparation).
The H2 and Br γ line emission are thus consistent with UV excitation, from η Car. The clumps themselves are swept-up cloud material, likely driven by the mechanical luminosity of a strong wind from η Car. This wind is strong enough to also produce shocked H2 emission. Since there is now no clear evidence for any such emission this wind cannot still be interacting with the clumps. The clumps must currently be coasting while being externally heated by the radiation field from the nearby stars. Considering the episodic events of η Car's evolution the current state of the clumps may only represent a quiescent stage in their life and they may be overrun in the future by another fast stellar wind from η Car, resulting in shocked emission. This type of variability however would be on time scales of the order of 1 000 years, considerably longer than the observed 5.54-year brightness fluctuation cycle of η Car (Damineli, Conti, & Lopes 1997).
It is possible to estimate the number of PAH molecules that are emitting from the surfaces of the clumps, using their 3.29 µm flux and the incident far-UV radiation field. Following Fig. 13 of Allamandola, Tielens, & Barker (1989), the typical energy emitted in the 3.3 µm band by a PAH molecule, excited via far-UV fluorescence to 30 000 cm −1 of vibrational energy, is ∼ 6 × 10 −14 erg s −1 molecule −1 sr −1 . For an incident far-UV field, G0 ∼ 10 4 , there are ∼ 10 12 photons cm −2 s −1 . Taking a typical absorption cross-section for a PAH molecule of 7 × 10 −18 cm 2 per C atom (Tielens, private communication), then for a typical 50-atom PAH  Figure 4. A schematic representation of the three dimensional geometry of the region, based on the data presented here. The region of ionized gas is bounded by an ellipse and the star denotes the location of η Car. Each clump, with the exception of U1, has been grouped according to its velocity and optical features (see Table 1). The arrows indicate the direction in which the clumps are moving with respect to the ionized gas.
molecule we expect an excitation rate of ∼ 4 × 10 −4 s −1 , or about one excitation each hour. This energy is rapidly re-emitted through infrared fluorescence in the various PAH emission bands, including the 3.3 µm band. For the peak 3.29 µm flux of 3 × 10 −3 erg s −1 cm −2 sr −1 measured from clump A1, we thus derive an emitting PAH column density of ∼ 10 14 cm −2 from this region. These estimates show that a significant fraction of the C-atoms may reside in PAH molecules. For instance, if the PAH emitting region were confined to an optical depth Av < 0.1 from the front surface of the cloud, it would imply a C-abundance in the PAH molecules of ∼ 3 × 10 −5 relative to hydrogen, or around 30 per cent of the total carbon, assuming typical interstellar carbon abundances.
If the rate of destruction of H2 molecules by far-UV radiation exceeds their formation rate on dust grains, it is possible that photodissociation equilibrium will not exist in the molecular clumps. For instance, Hollenbach & Natta (1995) derive a characteristic time-scale of 5 × 10 8 n −1 years to establish equilibrium. This translates to ∼ 5 × 10 4 years using the densities derived here, which is greater than the assumed age of ∼ 10 3 years for the wind. If this were the case the ionization-dissociation front impinging on the clumps will constantly be exposing fresh molecular material to the radiation field and the fluorescent line intensities will be higher than steady-state model predictions. However the front, moving at perhaps ∼ 1 km s −1 , would not have had time to propagate through the clumps. If photodissociation equilibrium does not exist we would expect the ratio of the ν=1-0/2-1 S(1) lines to preferentially show pure fluorescent values (∼ 2) where the ν=1-0 S(1) line is strongest (see Allen et al. 1999). This could be tested with further observations.

CONCLUSIONS
We have imaged the Keyhole region of the Carina Nebula in the H2 1-0 S(1) line at 2.122 µm and in the PAH emission feature at 3.29 µm, and compared the results to optical, radio continuum and 12 CO(2-1) images of the region. The H2 and PAH images are remarkably similar, indicating they are fluorescently excited, and show a series of emission clumps typically 30-60 arcsec in extent distributed around the edges of the radio emission. When compared to the optical and CO images three categories of clumps are apparent: (i) dark regions of optical obscuration with CO but no H2 or PAH emission (clumps 1, 2, and 3 which form the Keyhole feature), (ii) faint, optically dark clumps with bright rims and H2 and PAH emission (clumps A1, A2, 4 and 5) and (iii) faint optically dark clumps with H2 and PAH emission (clumps B and C).
The clumps are typically of mass 10 M⊙ and appear to have been swept up from ambient cloud material by a wind from η Car. They have been driven in all directions away from η Car, and must now be coasting at speeds of ∼ 10 km s −1 as there is no evidence of any shocked emission. The prominent dark clumps of the Keyhole are coming towards us and are outside the H II region. The bright-rimmed clumps A1, A2, 4 and 5 are within the H II region and thus are being heated by the radiation field. Clumps such as A1 are seen face-on whereas clumps 4 and 5 are seen edge-on. Finally the faint dark clumps B and C are on the far side of the H II region and have been driven away from η Car in the opposite direction to ourselves. The clumps comprise the last remnants of the ambient molecular cloud from which η Car formed.