LARGE-SCALE EXTENDED EMISSION AROUND THE HELIX...

LARGE-SCALE EXTENDED EMISSION AROUND THE HELIX NEBULA: DUST, MOLECULES, ATOMS,AND IONS

A. K. Speck, M. Meixner, D. Fong, P. R. McCullough,1D. E. Moser, and T. Ueta

Department of Astronomy, University of Illinois at Urbana-Champaign,MC-221, 1002West Green Street, Urbana, IL 61801Received 2001 August 28; accepted 2001 October 10

ABSTRACT

We present new observations of the ionized gas, molecular gas, and cool dust in the Helix Nebula (NGC7293). The ionized gas is observed in the form of an H� image, which is constructed using images from theSouthern H� Sky Survey Atlas. The molecular emission was mapped using the H2 v = 1 ! 0 S(1) line at2.122 lm. The far-infrared (FIR) observations were obtained using ISOPHOT on the Infrared Space Observ-atory. The H� observations are more sensitive than previous measurements and show the huge extent of theHelix, confirming it as a density-bounded nebula and showing previously unseen point-symmetric structures.The H2 observations show that the molecular gas follows the distribution of molecular material shown in pre-vious work. The molecular emission is confined to that part of the nebula seen in the classic optical image.Furthermore, comparison of the H2 emission strength with time-dependent models for photodissociationregions (PDRs) shows that the emission arises from thermal excitation of the hydrogen molecules in PDRsand not from shocks. The FIR observations, at 90 and 160 lm, represent mostly contributions from thermaldust emission from cool dust grains but include a small contribution from ionized atomic lines. Comparisonof the FIR emission with the H� observation shows that the dust and ionized gas are coincident and extendto �110000 radius. This equates to a spatial radial extent of more than 1 pc (assuming a distance to the Helixof �200 pc). Assuming that the outer layers of the circumstellar shell have spherical symmetry, radiativetransfer modeling of the emission in H� gives a shell mass of�1.5M�. However, the modeling does not coverthe outermost part of the shell (beyond �60000 radius), and therefore this is a lower limit for the shell mass.Moreover, the models suggest the need for very large dust grains, with�80% of the dust mass in grains largerthan 3.5 lm. Comparison of these new observations with previous observations shows the large-scale stratifi-cation of the Helix in terms of ionized gas and dust, as well as the coexistence of molecular species inside theionized zones, where molecules survive in dense condensations and cometary knots.

Key words: circumstellar matter — planetary nebulae: individual (NGC 7293) —stars: AGB and post-AGB — stars: evolution — stars: mass loss

1. INTRODUCTION

At a distance of only �200 � 50 pc (Harris et al. 1996;Harrington & Dahn 1980), the Helix Nebula has a largeangular extent on the sky (�140 radius at the largestextent; O’Dell 1998; Malin 1982), making it extremelyvaluable for the study of the spatial distributions of thedifferent species in nebulae (i.e., ions, atoms, molecules, anddust). Understanding the distributions of the differentspecies enables a detailed analysis of the physical and chemi-cal conditions during the postasymptotic giant branch(post-AGB) phase and allows a better understanding of thechemical enrichment of the interstellar medium (ISM) fromthese intermediate-mass stellar sources.

The current work looks at the large-scale structure anddistribution of species across the nebula. We present newobservations at both near- and far-infrared (IR) wave-lengths, showing the distribution of the molecular hydrogenand cool dust grains, respectively, as well as a new deep H�image, which shows that the ionized part of the nebula has alarger extent than previously seen. These new observationsare compared with those at other wavelengths from radio toX-ray to obtain a comprehensive picture for the large-scalestructure of the Helix.

The large-scale structure of the extended emission aroundthe Helix has been widely studied for the last fifty years. Thecurrent paradigm for the structure of the nebula is that of adisk composed of concentric rings in which the level of ion-ization decreases with distance from the central star. Thecentral zone, which appears to be a hole in the classic opticalimage, is characterized by He ii emission, from recombina-tions of He2+ (O’Dell 1998), and [O iv] emission (Leene &Pottasch 1987). Outside this central, highly ionized regionthere is a ring of lower ionization species, that coincideswith the inner part of the optical ring and is characterizedby [O iii] and H� emission (Warner & Rubin 1975; O’Dell1998; Henry, Kwitter, & Dufour 1999). The H� emissionpersists beyond the [O iii] ring and also coincides with theouter part of the optical ring characterized by the even lowerionization species [O ii] and [N ii] (O’Dell 1998; Henry et al.1999). To date most optical images do not show ionizedemission beyond the classic nebula. However, O’Dell (1998)and Malin (1982) both show an arc of ionized emission outat an angular radius of�85000.

Observations of molecular species, CO (Huggins & Healy1986, 1989; Healy & Huggins 1990; Forveille & Huggins1991; Huggins et al. 1992; Young, Phillips, & Knapp 1993,hereafter YPK) and H2 (Storey 1984; Kastner et al. 1996;Cox et al. 1998), show that the molecular emission is coinci-dent with and does not extend beyond the ionized emission.Molecules are expected to be destroyed by the passage ofthe ionization front and excited in the photodissociation1 Cottrell Scholar of Research Corporation.

The Astronomical Journal, 123:346–361, 2002 January

# 2002. The American Astronomical Society. All rights reserved. Printed in U.S.A.

346

region (PDR) that exists at the boundary between the ion-ization front and the molecular gas. For this reason, it isexpected that there should be a ring of neutral and molecu-lar species outside the rings of ionized gas. This is not seenin the observations. The coincidence of molecular emissionwith the ionized nebula implies that these species are some-how shielded from the ionizing Lyman continuum by beinginside dense (optically thick) condensations of material.Indeed, evidence for such condensations has been knownfor nearly fifty years (e.g., Zanstra 1955; Aller 1954), andthey are now known as the ‘‘ cometary knot’’ (e.g., O’Dell &Handron 1996 and references therein). Dust exists through-out the nebula, as is shown in the IRAS far-IR images pre-sented by Leene & Pottasch (1987).

In x 2 we present our new observations of the spatialdistribution of ionized gas (H�), molecular hydrogen,and cool dust. Section 3 compares the new observationswith the numerous studies at other wavelengths and dis-

cusses qualitatively the implications of the distributionsof the various species in terms of the large-scale struc-ture of the nebula. In x 4 we show the results of model-ing the nebula by using a simple spherically symmetricradiative transfer model. The conclusions are discussedin x 5.

2. OBSERVATIONS

We have observed the Helix in H� ionized emission, far-infrared (FIR) cool dust continuum emission and near-IR(NIR) molecular hydrogen lines.

2.1. H�Observations

The continuum-subtracted, deconvolved H� image of theHelix, presented in Figure 1, was constructed from imagesfrom the Southern H-Alpha Sky Survey Atlas (SHASSA;Gaustad et al. 2001). We aligned and averaged four of the

Fig. 1.—SHASSA H� observation of the Helix Nebula. The lowest contour is the 1 � level of 3 � 10�7 ergs cm�2 s�1 sr�1. Higher contours correspond tologarithmic brightnesses of �6.3, �6.0, �5.7, �5.4, �5.1, �4.8, �4.5, �4.2, �3.9, and �3.6. Point-symmetric features in the northeast-southwest, east-west,and southeast-northwest directions are marked as A1–A2, B1–B2, and C1–C2, respectively.

EMISSION AROUND THE HELIX NEBULA 347

SHASSA images whose IDs end in ‘‘ .f.’’ into one image.This image showed some faint loops extending to r = 200

from the Helix’s center, superposed on a faint, azimuthallysymmetric halo. This latter halo was removed by five itera-tions of a Lucy-Richardson deconvolution, which allowed abetter look at the point-symmetric structures extending outfrom the Helix as far as�180 (see alsoMalin 1982).

The central ringlike structure in Figure 1 is the classicoptical ring of the Helix. The most remarkable aspects ofthis new deep H� image are the spatial extent of the emis-sion and its point-symmetric structure. The ionized gasassociated with the Helix Nebula has an angular radius of�110000 from the central star. At a distance of 200 pc thisgives the Helix a physical radius of�1.1 pc. This image alsoshows that the Helix is a density-bounded nebula (i.e., theionizing photons from the central star reach beyond theedge of the material that makes up the nebula), with the neb-ular H� emission merging with the background ISM H�emission. The shape of the extended ionized emission iselongated in the east-west direction and shows evidence ofpoint-symmetric structures in the northeast-southwest,east-west, and southeast-northwest directions, marked inFigure 1 as A1–A2, B1–B2, and C1–C2, respectively. Thearc of ionized emission previously seen by O’Dell (1998) andMalin (1982) can easily be seen to the northeast of the mainring, within the extent of the point-symmetric structures.

To determine how much of the azimuthally symmetrichalo was due to the wings of the point-spread function(PSF), we created an azimuthally symmetric PSF from hun-dreds of stars in scores of SHASSA images and deconvolvedthe image using this PSF. The PSF of a very bright star suchas Sirius appears to be azimuthally symmetric (aside fromhuge bleed trails), so we have assumed the PSF is azimu-thally symmetric to get adequate signal for the PSF toapproximately r = 400. The profile of the PSF is shown inFigure 2, together with profiles of both the original anddeconvolved H� images of the Helix.

We prepared the image for the Lucy-Richardson decon-volution method by converting to photon units, forcing neg-ative pixels (residuals around stars from continuumsubtraction) to zero, and adding in a DC level from telluricemission lines. (The DC level was removed after the decon-volution). For the contour plot in Figure 1, we smoothedthe deconvolved image with a 3 pixel by 3 pixel medianfilter.

2.2. ISOPHOT FIRObservations

The FIR observations were obtained using the InfraredSpace Observatory (ISO). We have obtained FIR linearscans of the Helix Nebula using ISOPHOT (the imagingspectro-photopolarimeter on ISO; Lemke et al. 1996).These linear scans, which are centered on the central star ofthe nebula, represent traverses across the entire circumstel-lar dust shells. The scans were taken on 1997 May 5 at aposition angle of 155� east of north. The position, orienta-tion and extent of the linear scans relative to the opticalimage of the Helix is shown in Figure 3. Both forward andreverse scans along this track were obtained to determinethe repeatability of structures observed in the object’s spa-tial profile. The scans were imaged with two filters, the C10090 lm, which has a 3 � 3 pixel format with 4600 pixel scale,and the C200 160 lm, which has a 2 � 2 pixel format with9200 pixel scale. The PHT32 AOT uses a combination of ras-ter mapping and chopper sweeping to create a map. Foreach raster step (of 6000 for the 90 lm and 9200 for the 160 lmfilter), the chopper was used to make smaller steps of 1500

and 3000 to give an image pixel size of 1500 � 4600 and3000 � 9200 for the 90 and 160 lmfilters, respectively. The lin-ear scans used the maximum raster length possible of 300

and 460 for the 90 and 160 lm filters, respectively. The chop-per, however, sweeps across the raster position, whichincreases the total length of the linear scans to 360 for the 90lm and 530 for the 160 lm filter. The point-spread functions(PSFs) have full-width half-maxima (FWHM) of 44>5 forthe 90 lm filter and 97>4 for the 160 lm filter. The ISO-PHOT data presented in this paper were reduced using thePHOT Interactive Analysis package (Gabriel et al. 1997),together with the P32Tools package developed specificallyfor reduction of the AOT PHT32 data (Schulz & Peschke2002). The P32Tools package allows a better calibration ofthe PHT32 data, so that the errors on the calibrated datacan be reduced to�30%.

The flux-calibrated profiles of the ISO linear scans areshown in Figure 4. We also show the ISO PHT22 data fromthe ISO archive (principal investigator, Cox) for compari-son. The PHT22 observations have a lower angular resolu-tion and a smaller image area compared with our PHT32observations. Slices through the PHT22 images at the sameposition angle as the PHT32 linear scans are also includedin Figure 4. The errors on the photometry of the ISOPHOTdata are�30%. Therefore, the two 90 lm scans (PHT22 andPHT32), shown in Figure 4 (top), are consistent within thecalibration errors. The PHT32 160 lm and PHT22 180 lmscans show very similar morphologies. In fact the main dif-ference between the 160 and 180 lm filters is that the 160 lmfilter is 20 lm broader to the short-wavelength end (seeTable 1), which accounts for the difference in brightness ofthe nebula seen by these two filters. The 160 lm scan showsevidence of low-level dust emission all the way out to anangular radius of �110000 from the central star. This is also

Fig. 2.—Slice through the H� image, together with the point-spreadfunction (PSF) of the H� image, as determined from images of stars andtwo PNs (NGC 7009 and NGC 6572). The two PNs have the same radialprofile as stars, even though the former are pure line emitters, whereas thelatter are continuum emitters. Thus, we can confidently use a PSF deter-mined from many stars to deconvolve the line emission of the Helix. Theazimuthally symmetric PSF that we adopted is the dashed line; it has aGaussian core with FWHM = 1<3, an r�2 region out to 110, and an r�3 pro-file beyond that. The Gaussian core-plus-r�2 profile has been observed inmuch larger Schmidt optics, albeit with different scale factors (King 1971).The r�3 outer profile avoids the logarithmic divergence of the integral toinfinity of an r�2 profile.

348 SPECK ET AL. Vol. 123

seen in the IRAS data (see Fig. 5). Unfortunately the otherdata (PHT22 90 and 180 lm, and PHT32 90 lm) do notcover as large a distance from the central star, and thereforethe behavior of the outlying dust emission at these wave-lengths is difficult to determine.

2.2.1. Comparison with the IRAS Images

Figure 5 shows the ISO PHT32 observations togetherwith the ISO PHT22 and IRAS 60 and 100 lm data. It isclear that the ISO PHT22 and ISO PHT32 90 lm profilesare very similar to the IRAS 60 and 100 lm profiles. Themorphology of the ISO PHT32 160 lm and the ISO PHT22180 lm profiles, while similar to one another, show amarked difference in comparison with the other data. Theyare clearly missing the central region emission seen in theshorter wavelength profiles (at 60, 90, and 100 lm). This isdiscussed further in x 3.1.

The motivation for observing the Helix at the long wave-lengths available to ISOPHOT was to determine the distri-bution of cool dust associated with the nebula. Before wecan properly compare the FIR images at different wave-lengths we need to understand the contribution to the emis-sion seen in these broadband images from the narrowemission lines that are characteristic of planetary nebulae(PNs). To do this, we have investigated the ISO LWS01

full-grating observations (43–196 lm; Clegg et al. 1996)available from the ISO data archive (principal investigator,Cox). The LWS OLP9.5–processed data were analyzed fol-lowing standard reduction procedures by using ISAP 2.0a.The subspectra were merged to form a continuous spectrumby comparing the continuum levels in the overlappingregions. Line fluxes were measured by fitting a Gaussian tothe line profile.

We have PHT22 and PHT32 scans using three differentfilters (90, 160, and 180 lm), all of which are potentially con-taminated by line emission. Table 1 shows the emission linesthat fall within the filters of both the IRAS and the ISOobservations. The 90 lm filter covers the wavelength range69–121 lm, which includes the [O iii] 88 lm fine-structureline. The 160 lm filter covers the 129–219 lm range, and the180 lm filter covers the 150–221 lm range, both of whichinclude the [C ii] 158 lm and the [N ii] 205 lm fine-structurelines. Comparing the flux contained in the emission lines tothe flux in the continuum for these broadband filters, wefound that the [O iii] 88 lm line contributes less than �10%of the flux in the 90 lm band. For the 160 and 180 lm bandswe can obtain only rough estimates, since the LWS01 spec-trum goes only to 196 lm, while these long-wavelength fil-ters go out to �220 lm. However, we still find that the [C ii]158 lm emission line contributes less than �20% to theemission in these bands.

The broadband FIR images obtained by IRAS are alsopotentially contaminated by nebular line emission. The 60lm band covers the 30–84 lm wavelength range, and the100 lm band covers the 70–140 lm range. The IRAS 100lm image is very similar to that of the ISOPHOT 90 lm

Fig. 3.—Extent, position and orientation of the ISO linear scans withrespect to the classic optical image of the Helix

Fig. 4.—Profiles of the ISOPHOT PHT32 linear scans of the Helix,together with slices through the lower resolution two-dimensional ISO-PHOT PHT22 images taken at the same position angle as the linear scans.The dotted lines represent the point-spread functions.

No. 1, 2002 EMISSION AROUND THE HELIX NEBULA 349

images, as one would expect. The 60 lm image is not conta-minated by the 88 lm [O iii] fine-structure line but is poten-tially contaminated by the [O iii] 52 lm and the [N iii] 57 lmfine-structure emission lines, as well as the [S iii] 33 lm line.

We concur with Leene & Pottasch (1987) that the ionizedemission lines do not contribute significantly to the emissionin these broad FIR bands, with the majority of emission inour observations coming from cool dust. The LWS spectrawere taken at the position of the optical ring of the Helixand not of the central region; therefore these estimates holdonly for the ring region. The optical ring region of the Helixis bright at all FIR wavelengths. However, the ISO LWSspectra of the optical ring region do not show any signaturedust features or even a peak in the continuum suggesting arange in dust grain sizes. Therefore, a particularly distinctgrain size population is needed to produce such a spectrum.Further radiative transfer modeling is required to constrainthe grain size distribution.

2.3. Molecular Hydrogen Observations

To make the NIR H2 map of the Helix we used the Near-Infrared Imager (NIRIM; Meixner, Young Owl, & Leach1999) at the Mount Laguna 1 m telescope2 on 2000 August2–5. The imager was used with the 200 pixel scale, giving the256 � 256 pixel array an imaging area of 8<5 � 8<5. Thelarge angular size of the Helix Nebula demanded mosaic

mapping. The mosaic was made from images taken at 13positions in a diamond shape centered at the position of thecentral star at (22h29m38 95, �20�50013>5). Positions fartherfrom the central point were observed, but the entire molecu-lar emission from the nebula was found to be containedwithin the 13 position diamond mosaic whose maximumextent was 12<25 from the central star position 2h29m38 95,�20�50013>5. Off positions were taken regularly at offsets of0=5 from the central star. The mosaic map was imaged atthree wavelengths: the v = 1 ! 0 S(1) line of H2 at 2.122lm, Br � at 2.166 lm, and the broadband K0 at 2.12 lm.Images taken with theK0 filter had a typical integration timeof 3 s, while the images taken using the H2 and Br � filtershad typical integration times of 80 s. The data were reducedand compiled using IRAF. Each exposure was flat-fieldedto eliminate large pixel-to-pixel sensitivity variations in thedetector array and was then sky-subtracted. Sky emissionmaps were constructed by taking the average of several skyframes obtained close to the target field both in time andspace. We used point sources in each field as referencepoints to determine relative offsets between frames. Finallythe frames were co-added using the derived offsets. Theresulting H2 image is shown in Figure 6. The observationswere calibrated using standard stars: BS 8709, PLX 5546,HD 162208, and HD 203856. These stars were observedusing the H2, K

0, and Br � filters several times throughouteach night to account for changing sky conditions. The K-

Fig. 5.—Comparison of all FIR data. In each case for which a two-dimensional image is available (i.e., the PHT22 data and the IRAS 60 and100 lm images), the profile represents a slice through the image at the sameposition as our ISOPHOT PHT32 linear scans. For each data set an offsethas been added for clarity: the IRAS 60 lm data have an offset of 36 MJysr�1; IRAS 100 lm data have an offset of 12 MJy sr�1; ISO PHT22 90 lmdata have an offset of 7 MJy sr�1; ISO PHT32 90 lm data have an offset of8MJy sr�1; and the ISO 160 and 180 lmdata are not offset.

TABLE 1

Emission Lines within the ISO IRASBroad FIR Filters

Major Lines Minor Lines

IRAS 60 lm, 30–84 lmRange

[S iii] 33 lm ......... [Si ii] 35 lm

[O iii] 52 lm......... [Ne iii] 36 lm

[S i] 56 lm ........... [F iv] 44 lm

[N iii] 57 lm ........ [Fe iii] 52 lm

[O i] 63 lm........... [P ii] 61 lm

[F ii] 67 lm

[Si i] 68 lm

ISO 90 lm, 69–121 lmRange

[O iii] 88 lm......... [Al i] 89 lm

[Fe iii] 105 lm

IRAS 100 lm, 70–140 lmRange

[O iii] 88 lm......... [Al i] 89 lm

[N ii] 122 lm........ [Fe iii] 105 lm

[Si i] 130 lm


[O i] 146 lm......... [Si i] 130 lm

[C ii] 158 lm........

[N ii] 205 lm........


[C ii] 158 lm........

[N ii] 205 lm........

2 Mount Laguna Observatory is jointly operated by San Diego StateUniversity andUniversity of Illinois at Urbana-Champaign.

350 SPECK ET AL.

Fig. 6.—Helix Nebula imaged in H2 2.122 lm v = 1 ! 0 S(1) line

band fluxes for these standard stars were obtained from theElias standard-star list for HD 162208 and HD 203856(Bouchet, Manfroid, & Schmider 1991 and Carter 1990 forBS 8709; Veeder 1974 for PLX 5546). We did not correct forthe small changes in air mass because the standard stars BS8709 and PLX 5546 are close to the Helix and the timebetween consecutive observations was small.

The flux calibration gives a peak value for the H2 emissionof �3 � 10�4 ergs s�1 cm�2 sr�1. The average brightness is�2 � 10�4 ergs s�1 cm�2 sr�1. Our Br � images show thatthe Helix was undetected in this line, giving a 1 � upper limitto the Br � emission brightness of �7 � 10�8 ergs s�1 cm�2

sr�1.TheK0 images are consistent with all the nebular emission

in this filter coming from the H2 v = 1 ! 0 S(1) 2.122 lmline. The K0 images also provide a 3 � detection of the cen-tral star in this band. This shows that the flux of the centralstar at K0 (2.12 lm) is 2 � 10�13 ergs s�1 cm�2 lm�1. Thecentral star of the Helix is believed to be a 123,000 K whitedwarf (Bohlin, Harrington, & Stecher 1982). Using the V-band magnitude for the Helix Nebula’s central star quotedin Bohlin et al. (1982) to normalize a 123,000 K blackbody,we found that the expected K0 flux density is �7 � 10�13

ergs s�1 cm�2 lm�1. Therefore theK0 flux we detect from thecentral star is less than that expected from the 123,000 Kwhite dwarf and does not show evidence for an M dwarfcompanion, as suggested by Guerrero et al. (2001) toaccount for the hard X-ray source coincident with theHelix’s central star.

3. STRUCTURE OF THE HELIX

Comparison of our new observations, together with pre-vious work, shows that the Helix is composed of three dis-tinct zones that we will discuss in the following sections.This structure is shown as a schematic cartoon in Figure 7.

3.1. Central Zone

The central zone is characterized by emission from highlyionized gas (He ii and [O iv]), as well as some odd emissionat 60–100 lm FIR wavelengths, and appears to be a hole inthe classic optical image. This zone has a spatial extent of�25000, with its outer edge bounded by cometary knots.There is no evidence of molecular emission within a radiusof�9000 (e.g., O’Dell, Henney, & Burkert 2000). This zone iscompletely filled with emission from He2+ (O’Dell 1998)and O3+ (Leene & Pottasch 1987). Our observations showthat there is also H� emission, albeit at a lower level than inthe ring (see also Henry et al. 1999; O’Dell 1998). The mostextraordinary observations of the central region are the FIRdata (IRAS and ISO). These show that there is significantemission in the central zone at 60, 90, and 100 lm that is notseen at 160 or 180 lm (see x 2.2) or at mid-IR wavelengths(Cox et al. 1998).

Since the ISO spectra suggest that there is negligible con-tribution from the [C ii] 158 lm and [N ii] 205 lm emissionlines to the 160 and 180 lm scans, we can assume that theemission in these bands arises solely from dust. Dust thatemits at 160 and 180 lm will also contribute to the contin-uum emission in the 60–100 lm bands. Let us also assumethat the same dust component is also the major contributorto the flux in the IRAS 60 and 100 lm bands and the ISO 90lm band. Therefore the excess emission in the central region

of the 90 lm scan, as seen in Figure 8, is not from the samedust responsible for the emission coincident with the opticalring at about�25000.

To understand the difference between these two scans the160 lm scan has been normalized to the 90 lm scan. Thenormalized 160 lm scan is shown together with the 90 lmscan in Figure 8 (top). In Figure 8 (bottom) the 160 lm scanhas been subtracted from the 90 lm scan, so that the excesscentral emission at 90 lm is clear. Also shown are the IRAS25 lm [O iv] data (Leene & Pottasch 1987) and the He ii

data from O’Dell (1998). This shows that the 90 lm excess isspatially coincident with He2+ and O3+ ions and thereforewith the most ionized part of the nebula.

There are two possible explanations for this excess emis-sion in the central zone: dust or FIR fine-structure emissionlines. First, let us consider dust. The IRAS 60 lm image hasthe same morphology as the IRAS 100 lm image (Leene &Pottasch 1987) and the ISO 90 lm data. Therefore, if theexcess emission is due to dust, it must also contribute a simi-lar amount of emission in the 60 lm band as in the 90 and100 lm bands. Furthermore, Cox et al. (1998) found thatthe mid-IR (5–16 lm) emission from the Helix is almostexclusively due to H2 rotational lines and argon and neonemission lines, with no signature dust features and very littleunderlying continuum. Therefore, if dust is responsible forthe excess emission, a grain population must exist that emitsstrongly at both 60 and 90 lm, but not at either 15 or

Clumps

3+O

0He+He

2+

O

He

shadow

+

H

2+

2

CO

O

+

H+

H0H

Molecular

❂

lightstar

dust

dust

Fig. 7.—Schematic view of a density-bounded planetary nebula, whereboth the ionization front and the shock front have passed through the entirevisible nebula. The edge of the nebula is where the density is so low that wecan no longer see the emission. The knots provide many mini-ionizationfronts throughout the nebula. The inset shows a magnified view of a singleglobule with the ionization front and PDR around the edge. Note that dustis present everywhere there is gas.


160 lm. This narrow wavelength coverage for dust emissionseems peculiar. The nature of the implied dust grain size dis-tribution is discussed further in x 4.

It is also possible that the excess central emission at 60and 90 lm is caused by atomic or ionized emission lines.The emission lines that fall into the broadband filters ofthe ISO and IRAS images are listed in Table 1. The onlymajor atomic or ionized emission line that falls withinthe wavelength range of the 90 lm filter is the [O iii] 88lm line. This line also contaminates the IRAS 100 lmimages, while the other [O iii] fine-structure line at 52 lmfalls within the range of the 60 lm filter. Comparisonbetween our 90 lm scan and a slice through the [O iii]5007 A image from O’Dell (1998) taken at the same posi-tion angle as our ISO data shows that they have similarbulk morphologies (Fig. 9). However, the inner region ofthe 5007 A image does not show the extra emission wesee at 90 lm (Fig. 8, bottom). The 5007 A emission fromO2+ covers a larger angular extent than the 90 lm excessemission and peaks at the inner edge of the classic opticalring, while the 90 lm excess emission is strongest in thecentral region of the Helix. Therefore, if the excesscentral emission is to be attributed to [O iii] fine-structureemission lines, there must be a mechanism by which thespatial distribution from different lines of the same spe-cies can have markedly different morphologies. While it

may be possible to manipulate the relative intensities ofthe optical and FIR fine-structure lines by using thedifference in the critical densities for these lines, the resultwould be to increase the 5007 A line relative to the 88lm line, not vice versa. Therefore, it seems unlikely thatsuch huge differences in morphology as seen between the5007 A and 90 lm emission could be explained this way.Therefore, we believe that the excess emission in the cen-tral zone is due primarily to dust. Model calculations ofdust emission are discussed further in x 4.

3.2. Ring

The ring region is characterized by the classic opticalimage and shows emission from ionized gas (H�, [O iii],[N ii], etc.), dust, and molecular gas (H2 and CO) in clumps.Figure 9 compares the spatial distributions of emission fromdifferent species. Ionized gas, dust, and molecular gas allhave their strongest emission in the ring region, althoughthe exact distributions vary slightly. The [O iii] 5007 A emis-sion is stronger in the central zone than the other ionizedspecies and drops off more quickly to the outer part of thering. The dust, as seen in the ISOPHOT linear scans, followsthe same distribution as the ionized gas (H�). The dust andH� emission extends farther than that of the other speciesshown in Figure 9. By comparing the new H2 observations

Fig. 8.—FIR emission in the central region of the Helix: top, 90 lm scantogether with the normalized 160 lm scan; bottom, difference between the90 and normalized 160 lm linear scans compared with slices through theIRAS 25 lmdata and the He ii data fromO’Dell (1998) at the same positionangles as the ISOPHOT scan. The 25 lm and He ii data have been normal-ized to the ISO data. The He ii data have been convolved with a Gaussianto approximate the resolution of the ISO 90 lm scan. The dips below zeroon either side of the central peak are a result of the different angular resolu-tions of the 90 and 160 lmobservations.

Fig. 9.—Comparison of spatial distribution of emission from differentspecies: [O iii] 5007 A is from O’Dell (1998); [N ii] 6584 A and H� are fromHenry et al. (1999); H2 emission profiles are from NIRIM (this paper) andIRAS data; dust emission at 90 and 160 lm are from ISO (this paper).


with new deep H� observations (Fig. 10), it is clear that allthe molecular gas is entirely contained within the ionizedgas. This suggests that the molecular gas is shielded fromthe ionizing radiation by being inside denser (opticallythick) condensations. Indeed such condensations have beenidentified at the inner edge of the ring region and are knownas ‘‘ cometary knot’’ (see O’Dell & Handron 1996 and refer-ences therein).

3.2.1. Molecular Hydrogen Emission

H2 can be observed only where it is excited. H2 moleculescan be excited either collisionally, by shocks or high temper-atures (thermally), or radiatively, through fluorescence.Estimates of the H2 emission from PDR models led Cox etal. (1998) to suggest that there must be some (shocked) colli-sionally excited H2 emission from the Helix to account forits high intensity. However, PDRs in PNs are very differentfrom those associated with molecular clouds, to which mostpublished models refer. Most PDR models assume that anequilibrium point is reached at which H2 formation is equalto H2 destruction. However, for PNs, the PDRs are unlikelyto reach such an equilibrium since the timescale for H2 for-mation is generally long compared with the timescale for theturn-on of the far-ultraviolet (FUV) flux (Sternberg 1998).In this case the H2 emission may well be much higher thanexpected from equilibrium PDR models (e.g., Sternberg1998; Natta &Hollenbach 1998).

Natta & Hollenbach (1998) have modeled the emissionexpected from PNs, including the evolution of the centralstar and the ensuing PDR that propagates through the cir-cumstellar shell. In their models the shock that propagatesthrough and excites the molecular shell is due to the acceler-ation of the superwind to �25 km s�1 into the slower, lessdense AGB wind moving at �10 km s�1. They do notinclude the fast wind, since a shock produced by a fast wind(velocity�1000 km s�1) would dissociate the hydrogen mol-ecules rather than excite them. Since there is currently noevidence for a fast wind in the Helix (e.g., Patriarchi & Peri-notto 1991) its possible effect will not be considered here.Natta & Hollenbach (1998) also include the effect of X-raysproduced by hot central stars such as that of the Helix. Pre-vious models generally ignore X-rays and include only theFUV radiation. For a PN born of a high-mass progenitor,X-rays make a significant contribution to the PDR emis-sion, leading to much higher H2 emission than previouslyexpected. For their highest-mass progenitor (5 M�) andlowest-density nebula, the intensity of the H2 v = 1!0 S(1)line is�10�5–10�4 ergs s�1 cm�2 sr�1 after about 104 yr, con-sistent with but slightly lower than our observations, with apeak of�3 � 10�4 ergs s�1 cm�2 sr�1. The intensity is lowerfor lower mass progenitors and higher density nebulae.However, the progenitor mass for the Helix Nebula isexpected to be higher (6.5 M�; e.g., Henry et al. 1999;Gorny, Stasinska, & Tylenda 1997), which should yield ahigher flux from H2. Furthermore, while the models of

Fig. 10.—Comparison of the extent of the molecular hydrogen with that of the ionized hydrogen


Natta & Hollenbach (1998) do not include clumpiness, theysuggest that further enhancement in the PDRH2 emission isexpected as a result of the photoevaporation of the comet-ary knots and filamentary condensations seen in the Helix.This process effectively causes rapid advection of H2 fromwithin the condensations out into the mini-PDRs at theknot surfaces, which can lead to significantly higher H2

intensities than otherwise predicted.Therefore it is clear that previous ‘‘ static ’’ models of

PDRs are inadequate for the case of PNs in general and par-ticularly for clumpy nebulae such as the Helix. Based on themodels of Natta &Hollenbach (1998) we argue that the highlevel of H2 emission seen in the Helix is consistent with thatexpected from the PDRs in an evolved PN, that consists ofthe mini-PDRs that occur as a result of the clumpy natureof the molecular shell.

3.2.2. Extent of theMolecular Emission

Since all the molecular emission appears to come fromcometary knots and related condensations, what does thistell us about the evolution of this PN? Various formationmechanisms for these knots have been put forward, startingwith Zanstra’s suggestion that they formed from clumps ofmaterial already present in interstellar space. Since then theformationmechanisms have fallen into two basic categories:(1) clumps exist prior to the PN phase and the knots remainafter the passing of the ionization front and/or the fast windshock front (e.g., Zanstra 1955; Dyson et al. 1989) and (2)the knots form after the onset of the PN phase, arising fromRayleigh-Taylor (R-T) instabilities at either the ionizationfront or the fast wind shock front (e.g., Mathews 1968; Cap-riotti 1971, 1973). At present there is no conclusive evidenceto support either mechanism outright. Whatever their for-mation mechanism may be, these dense condensations areimportant as they are the reason we can still observe molec-ular emission from an evolved planetary nebula.

The inner edge of the optical nebula, at �10000, where theinnermost cometary knots lie, marks how far the circumstel-lar shell that resulted from AGB mass loss has drifted sincethe end of the AGB mass-loss phase. Assuming that the cir-cumstellar shell simply expands with constant velocity awayfrom the star, this implies that the AGB dust shell hasexpanded by �10000 in every direction since the end of theAGB. The molecular emission, as seen in both H2 and CO,currently extends to �40000–50000 radius. Therefore, at theend of the AGB the molecular shell would have extendedout to�30000–40000. At a distance of�200 pc, this equates toa radial extent for the molecular envelope at the end of theAGB of 0.3–0.4 pc (�1018 cm).

According to Mamon, Glassgold, & Huggins (1988) theextent of the molecular envelopes around cool, evolved stars(such as AGB stars) is governed by photodissociation by theinterstellar radiation field (ISRF) and the ability of the cir-cumstellar shell to shield itself from that radiation. Theyfound that, depending on the AGB mass-loss rate, theextent of the CO envelope around AGB stars is expected tobe in the range 1016–1018cm. The upper limit is very similarto the inferred extent of the molecular envelope of the Helixat the end of the AGB.

However, the current molecular emission is easily seen tooriginate from clumps that reside inside the ionized nebula,which have allowed the survival of molecular species as theionization front passed through the nebula. If these clumpsare a result of the general AGB mass-loss processes, then

such clumps would also shield the molecular species fromdissociation by the ISRF, and themolecular envelope wouldsurvive to a much larger radius than predicted by the simplemodels of Mamon et al. (1988). Why, then, do we not seemolecular emission at a farther distance from the centralstar?

There are three possible scenarios: (1) If the molecularclumps are not self-gravitating they will gradually expand asthey drift away from the star and thus lose their ability toshield themselves from the ionizing radiation. In this waythe clumps would be confined to a region relatively close tothe central star. (2) The molecular material does not formclumps during the AGB phase and thus the extent of themolecular shell is limited by the ISRF photodissociation asmodeled by Mamon et al. (1998). In this case, the clumpynature of the molecular emission as we now see it is due toR-T instabilities caused by the passage of the shock front orthe ionization front through the relatively smooth molecu-lar envelope. Since the smooth molecular envelope was lim-ited to �1018 cm by the ISRF, this is also the extent of themolecular material as it forms into clumps. However, thisassumes that the clump formation occurred more or lessinstantaneously at the onset of the PN phase. (3) Theclumpy nature of the AGB mass loss manifests itself onlyduring the last stages of the AGB phase. Thus the earliermolecular envelope was smooth and subject to limitation bythe ISRF, while the later clumpier molecular material is allthat has survived both the attack of the ISRF and the pass-ing of the shock and ionization fronts with the onset of thePN phase. The evidence for such clumpiness in the AGBphase is seen in the maser spots, which do, indeed, occur inthe latest stages of AGBmass loss (see Dyson et al. 1989).

3.3. Outer Halo

The outer halo is characterized by low-level H� and FIRemission. Comparison of previous observations with ournewH� and FIR data shows that the classic optical image ismerely the tip of the iceberg, with the ionized gas and dustemission extending to a radius of �110000 (or �1.1 pc). Therelative distributions of the dust (as seen by the IRAS 100lm data) and the ionized gas (as seen in H�) are shown inFigure 11, which shows the contours of IRAS 100 lm data(processed through the HIRES algorithm to get better reso-lution; see Aumann, Fowler & Melnyk 1990) overlaid onthe H� emission image. Figure 12 shows the profiles of theemission in the ISO 160 lm and IRAS 60 lm and 100 lmbands, together with the H� emission profile. The profileshave been normalized to the ISO 160 lm emission profilefor better comparison. It is clear, in both Figures 12 and 13,that there is low-level emission in the FIR out to a radius of�110000. The direction of the ISO linear scans (155� east ofnorth) is such that the outer, low-level dust emission seen inFigure 12 coincides with the region of Figure 11 where thedust emission appears to extend beyond the ionized gasemission. However, the rapid drop-off with radius in ionizedgas emission relative to the dust emission is a result of theemission mechanism and not an indicator that there is dustwithout gas in these regions. Dust emission strength is givenby I� /

Rdl�� B�(T), where dl is the pathlength, � is

the density, �� is the emissivity of the dust grains, andB�(T) is the Planck function, while the H� emissionstrength is given by IH� /

Rn2e dl, where ne is the electron

density. Therefore, H� emission is more sensitive to the


decreasing density than the dust emission and drops offmore quickly. The dust emission is very sensitive totemperature. At the temperatures relevant to this outerhalo region, small changes in dust grain temperaturecan have a huge effect on the dust emission strength in theFIR. The difference in the distributions of the H� and theFIR dust emission can be explained by their differingdependence of density and temperature. It is also possiblethat the Helix has a structure like that of some PNs andproto–planetary nebulae (e.g., NGC 7027 and the Egg Neb-ula), with a bipolar structure superposed on spherical dustshells (Hrivnak, Kwok, & Su 2001; Sahai et al. 1998; Latteret al. 2000). The H� extended emission shows three point-symmetric pairs (marked in Figure 1 as A1–A2, B1–B2, andC1–C2), reminiscent of these younger objects, that may bedue to precessing jets (e.g., Miranda, Guerrero & Torrelles2001).

4. MODELING THE HELIX

4.1. Model

We have modeled the emission from various specieswithin the Helix using a simple radiative transfer model. Werefer the reader to Hoare (1990) for details concerning themodel. The radiative transfer assumes spherical symmetry,a blackbody of Teff = 123,000 K (Bohlin et al. 1982), and aconstant gas-to-dust ratio throughout the nebula. Themodel ionized gas emissions assume atomic abundancesfrom Henry et al. (1999). These abundances show that theHelix has an approximately solar carbon-to-oxygen ratioand is therefore oxygen-rich. For this reason the dust emis-sion is assumed to arise predominantly from silicates. Theoptical constants included in the model are for ‘‘ astronomi-cal silicate ’’ from Draine & Lee (1984, 1987). The dustgrains are in thermal equilibrium with the radiation field

Fig. 11.—Very extended 100 lm IRAS emission from the Helix compared with the H� emission


that includes direct stellar radiation, free-free continuumemissions, and line emissions (mostly Ly�).

This simple model has several limitations. First, themodel assumes spherical symmetry, which is obviously aproblem when applied to an object such as the Helix, whichis patently not spherical. Second, the model assumes thatthe gas and dust is smoothly distributed. Again, in the caseof the Helix, it is known that the nebula is not smooth, butrather that it contains clumps of denser material (e.g., thecometary knots). Finally, the largest grain radius that theprogram can include is 3.5 lm. If larger grains are present inthe Helix their effect must be inferred. However, having laidout the limitations and in the absence of something moreelaborate, using this simple model gives an indication of theeffect of the changing distributions of different ions, atoms,and dust grains on the expected emission profiles. In partic-ular it will help us to understand the very extended cool dustemission and its variation with wavelength (see x 2.2.1) interms of dust grain size distributions and gas-to-dust ratios.

The model assumes that the dust and gas are coupled, sothat the density distribution of the dust follows that of thegas. Therefore, the first step was to match the spatial distri-bution of the H� emission.

4.2. Gas Density Distribution

We have two sets of H� emission observations available.Our own, new data (see x 2.1) is deeper and covers a largerarea of sky than that of Henry et al. (1999) but with lowerresolution. Therefore both are valuable in this study of thespatial distribution of various species. Both data sets areshown in Figure 13. The difference in the breadth of thepeak of emission is due to the difference in resolution, whilethe difference in the extent of the observed emission is due tothe depth of the observations. Our data show that there islow-level H� emission extending to more than 30000 fartherthan that seen in the data of Henry et al. (1999).3

The input profile that gives the best fit to the spatial distri-bution of H� is shown in Figure 13, together with theobserved data. This profile is the result of �70 attempts atmatching the H� profile, starting from a simple r�2 densitydrop-off. The best-fit profile has very low density in the cen-tral region, which is constant at�12 cm�3 out to the edge ofthe optical ring. From the edge of the optical ring (at�10000)to the peak of the density distribution at 30000 from the cen-tral star, the density rises steeply, as r2. The peak in the den-sity profile is only �60 cm�3. From the peak position thedrop-off in density is inversely proportional to the radiusout to a radius of�50000, where the drop-off increases to r�4.The r�4 dependence was found empirically by trying to bal-ance having enough material to see the ionized emission butnot so much as to make the shell optically thick and thuscurtail the spatial extent of the H� emission. This rapiddrop-off (faster than r�2) is indicative of a superwind mass-loss episode in which the rapidly increasing mass loss leadsto such a distribution. It was not possible to find a densitydistribution that would fit the outlying low-level H� emis-sion beyond �60000. This is probably due to the sphericalgeometry of the model.

The low level of the central density (�12 cm�3) is muchlower than expected following the work of O’Dell (1998)

Fig. 12.—Very extended FIR emission from the Helix: profiles of theISO 160 lm (solid line), together with slices through the IRAS 60 and 100lm bands and the H� emission at the same position angles as the ISO-PHOT scan. The profiles have been normalized to the ISO 160 lm emissionprofile for better comparison.

Fig. 13.—Model emission profiles. Top, comparison of the model pro-files of dust emission at 90 and 160 lm with the observed ISO linear scanprofiles, showing 90 lmobserved data (asterisks) and 160 lmobserved data( plus signs). The 90 lm model profile has been normalized to the observa-tion to show the shape of the distribution. Middle, observed profiles repre-senting slices through the observed data taken at the same position angle asthe ISO linear scans. The solid line shows the best-fit model. Bottom, inputdensity profile.

3 In fact the low-level H� emission extends as much as 60000 farther thanthat seen byHenry et al. (1999), depending on direction; see Fig. 1.


and Henry et al. (1999) and leads to the He ii emission beingmore extended than has been observed (O’Dell 1998). Thisis probably due to the fact that our model assumes sphericalsymmetry. If the Helix were more appropriately modeled asa disk, it would be possible to include more material in thecentral zone. This would produce a nebula that is ionizationbounded in the plane of the disk but density bounded per-pendicular to the disk, giving rise to the confined He ii emis-sion seen in the disk but allowing for the very extended H�emission. This suggests that the extended H� emission is aprojection of a bipolar flow, with the optical ring as thewaist (see O’Dell 1998).

The Stromgren radius for hydrogen for our best-fit modelis beyond the edge of the shell, producing a density-boundednebula. Table 2 lists the important parameters for the best-fit model.

4.3. Modeling the Dust

Having fitted the shape of the H� emission reasonablywell, we concentrated on the effect of the grain size distribu-tion and gas-to-dust ratio on the emission at FIR wave-lengths where the majority of the emission comes from cooldust. These models are compared with the ISO observeddata.

For each set of model input parameters the grain size dis-tribution is described by ng = a�3.5, with the minimum(amin) and maximum (amax) grain sizes being specified foreach model. Figure 14 shows the effects of changing thegrain size limits. When only small grains (0.005–0.05 lm)are used, the emission very close to the central star is toohigh at all IR wavelengths. While high central emission isexpected at 60 and 90 lm, this central emission is notobserved for mid-IR wavelengths (e.g., Cox et al. 1998) orat 160 or 180 lm (see x 2.2.1). Removing the small grains(size distribution 1.0–3.5 lm) gives a much better fit to theobserved data in the central region. While there is still a littlemore central emission in the model profiles than observed,this is partly due to the projection of the spherical shape ofthe model onto the plane of the sky. However, in the case oflarge grains, the model profile falls far short of fitting theouter wings of the FIR emission. Unfortunately the modelis not capable of including different grain size distributionsin different regions. Changing the grain size distributionsimplies that there is a paucity of small grains in the centralregions for the dust emission strength to remain low in themid-IR but that small grains are needed in the outer parts ofthe optical ring and beyond to produce the high level of FIR

emission seen beyond �30000. The very low level emissionobserved in the ISO 160 lm and IRAS 100 lm data beyond�80000 is probably partly heated by the ISRF (see YPK; theEgg Nebula and CRL 618, Speck, Meixner, &Knapp 2000).It is also possible that the clumpy nature of the Helix plays arole here. The model assumes that the shell is smoothly dis-tributed, but the clumpy nature of the Helix may allow pho-tons to escape through the gaps between the dense clumpsand thus the ionized emission would not be as confined asthe model suggests. This could also contribute to the heatingof dust grains in the outer halo.

For the best-fit model the widest dust grain size distribu-tion was used (0.005–3.5 lm). However, the upper limit forthe grain size is set by the capabilities of the model. A muchlarger maximum grain size is necessary to account for theabsolute flux-calibrated brightnesses in the observed FIRdata. By increasing the maximum grain size and maintain-ing the gas-to-dust ratio (by mass), the number of smallergrains is decreased, and the mass from these grains goes intothe larger grains. Decreasing the number of small grainslowers the overall flux levels in the IR. To match theobserved brightnesses at 90 and 160 lm for the Helix byusing the 0.005–3.5 lm grain size distribution, the gas-to-dust ratio must be increased to greater than 1000 (muchhigher than the 100–200 usually assumed; see Fig. 15). The

Fig. 14.—Effect of the grain size distribution on the spatial profile of theinfrared emission. Top, 90 lm profiles. The asterisks show the profile of theISO 90 lm linear scan.Bottom, 160 lmprofiles. The asterisks show the pro-file of the ISO 160 lm linear scan. The model profiles are normalized to thering (�25000) in the observed emission profiles.

TABLE 2

Parameters for the Best-Fit Radiative

Transfer Model

Parameter Value

T*............................................. 123,000K

L*............................................. 100L�D.............................................. 200 pc

Grain size distribution, ng......... a�3.5

amin .......................................... 0.005 lm

amax.......................................... 3.5 lm

Peak density............................. 60 cm�3

Gas-to-dust ratio ..................... 1000

Mass in circumstellar shell........ 1.3M�


absolute value of the brightness is also decreased by remov-ing the small grains, but, as explained above, some smallgrains are necessary to match the spatial distribution of theFIR emission. Therefore, including very large grains (>3.5lm) has the effect of decreasing the number of small grainswithout changing the gas-to-dust ratio or the shape of theemission profile at these wavelengths. To determine whetherthere are very large dust grains in the Helix, observationsmust be made at even longer wavelengths (submillimeterand millimeter). Indeed, centimeter-sized grains have beenseen in proto–planetary nebulae (e.g., the Egg Nebula; Juraet al. 2000). If large grains are not present, this implies thatthe gas-to-dust ratio is very high and that there is much lessdust in the Helix than previously thought.

One odd aspect of the models is the behavior of the emis-sion within the central zone. One of our objectives withthese models was to understand the origin of the centralemission seen at 60 and 90 lm and not at longer (160 and180 lm) or much shorter (<25 lm) wavelengths. Figure 16shows that in the central zone the dust heating is dominated

by the central star, while the dust heating in the optical ringis dominated by Ly� photons. Therefore, the extra emissionin the central zone comes from heating of dust by the centralstar. However, the inclusion of material in the central regionmanifests itself as a point source centered on the central star.No dust density distribution could be found that wouldcause the excess central emission to appear as more than asimple central point source. As explained above, the centralzone needs to have a deficit of small grains relative to theouter parts of the optical ring. Careful manipulation of thegrain size distribution allows the appearance of the (pointsource) excess emission at 60 and 90 lm over and above thatseen at 25 and 160 lm. By restricting the grain sizes in thecentral zone to the large grains (>1 lm), the central zonemay show an observed excess at 60 and 90 lm withoutshowing extra emission at 25 and 160 lm. However, toentirely fill the central zone with such dust emission aparticularly distinctive size distribution that changes withradius must be invoked. Very close to the star, there can beonly very few small grains. The relative number of smallgrains needs to increase with distance from the star, out tothe edge of the optical ring. The increasing number of smallgrains may allow the excess emission to fill the central zone,rather than being confined to a point source. At the edgeof the optical ring the line heating, dominated by Ly�,begins to dominate the dust heating so that the dusttemperature increases again and the peak of the dustemission is coincident with the peak of the H� emission.

What mechanism can be responsible for the strange dustgrain size distribution necessary to produce the excess 60and 90 lm emission? Since the relative number of smallgrains needs to increase with distance from the central stars(at least as far as the edge of the optical ring), we proposethe following mechanism: The high-energy photons pro-duced by the hot central star have sufficient energy todestroy small grains. With increasing distance from the starthe destruction of small grains by high-energy photonsdecreases, producing the necessary grain size distribution.

4.4. Mass of the Nebula

Based on the model, the circumstellar shell mass for theHelix is �1.3M�. The model requires a gas-to-dust ratio of1000, but even with a more usual gas-to-dust ratio of 100,the contribution to the nebula mass from the dust is negli-gible. The mass of the molecular gas is 0.025 M�, and themass of neutral gas is 0.18 M� (Young et al. 1999), giving atotal nebular mass of �1.5 M�. However, the modelaccounts for only the inner �50000–60000 radial extent. Thelow-level emission out to more than 110000 will contributesignificantly to the shell mass, making this a lower limit.

5. CONCLUSIONS

We have presented new observations of the ionized gas(H�), cool dust (at 90 and 160 lm), and molecular gas (H2)in the Helix Nebula. The H� observations go deeper thanprevious measurements, revealing the huge extent of theHelix Nebula, as well as confirming it as (at least partially)density bounded. The H� emission extends to a maximumangular radial extent of �110000, where it merges with thebackground ISM H� emission. At a distance of �200 pc,this equates to a spatial extent for the ionized gas of 1.1 pc.

Fig. 16.—Relative contributions of the direct stellar radiation and theH� to the dust heating. These relative contributions are virtually identicalfor all grain sizes within the distribution. The contribution of the diffusescattered starlight is always negligible.

Fig. 15.—Effect of the gas-to-dust ratio on the absolute brightness of theinfrared emission. The asterisks show the flux-calibrated profile of the ISO160 lm linear scan.


Furthermore, there is evidence of point-symmetric struc-tures in the extendedH� emission.

The ISOPHOT FIR observations have contributionsfrom both ionized atomic lines and thermal emission fromcool dust grains. We have shown that the contribution tothe broadband FIR images from the ionized lines is �10%,with the majority of emission arising from cool dust grains.Furthermore, the FIR emission from the central zone sug-gests that there is a population of relatively large dust grainsand a paucity of small grains within the central region of theHelix. Moreover, to fit the shape of the observed emission,the grain size distribution must include an increasing pro-portion of smaller grains with increasing distance from thestar. This implies that high-energy photons from the hotcentral star have destroyed small dust grains close to thestar. Small grains are needed in the outer parts of the dustshell to account for the shape and extent of the FIR emis-sion at larger distances from the star. To match the flux lev-els of the FIR emission, our models needed gas-to-dustratios of �1000. This is much larger than the value �100–200 usually assumed. One possible explanation is that muchof the dust mass is hidden in very large dust grains (>3.5lm). Therefore, our results imply either that �80%–90% ofthe mass of the dust is contained in dust grains larger than3.5 lm in radius or that there is much less dust in the Helixthan previously thought. The cool dust emission has beenshown to extend to a radius of �110000. At the distance ofthe Helix this equates to a spatial radial extent of over 1 pc.Our models of the Helix account for emission out to only aradius of �50000–60000. There are two possible explanationsfor the observed outer emission, which is not seen in themodels. At such a great distance from the central star, theouter part of the dust shell may be partly heated by theISRF (see Speck et al. 2000; YPK). This additional heatingsource is not included in our model. Alternatively, themodel assumes that the shell is smoothly distributed (i.e.,not clumpy), whereas the Helix is obviously clumpy. It ispossible that the photons currently confined within the clas-sic ring region may be able to reach greater radii by escapingthrough the gaps between clumps. This would provide someextra heating to account for the more extended emission.

The H2 observations show that the molecular gas isentirely contained within the ring region of the Helix, coinci-dent with the bright ionized gas emission. Therefore all themolecular gas is contained within clumps and cometaryknots in the ring region. The origin of these clumps is stillnot clear. Comparison of the H2 emission strength withtime-dependent models for PDRs (Natta & Hollenbach1998) shows that the excitation of the H2 molecules isconsistent with such models, especially if the clumpiness ofthe molecular medium is included. It is not necessary toinvoke shocks to achieve the strong molecular emissionobserved.

Based on a spherically symmetric model for the Helix weget a circumstellar shell mass of �1.5 M�. However, themodel accounts for only the inner�50000–60000 radial extent.The low-level emission out to greater than 110000 will con-tribute significantly to the shell mass, making this a lowerlimit.

We are very grateful to Bob O’Dell, both for the use of hisdata (from O’Dell 1998) and for constructive conversations,also to R. Dufour and R. Henry for the use of their data(from Henry et al. 1999). We would like to thank the ISO-PHOT team (Carlos Gabriel, Rene Laureijs, SybillePeschke, and Bernard Schulz) at Vilspa, Spain, for theirhelp and patience with understanding the data reductionand calibration techniques. We are also grateful to MikeBarlow for numerous constructive conversations orcorrespondences, and to Elric Whittington for readingthrough the manuscript. The H� data are from theSouthern H-Alpha Sky Survey Atlas, which was producedwith support from the National Science Foundation.A. K. S. was supported by NASA JPL 961504 and NASASTI 7898.02-96A. M. M., T. U., and D. E. M. weresupported by NSF CAREER award AST 97-33697.D. F. was supported by The Laboratory for AstronomicalImaging at the University of Illinois and NSF grant 99-81363. P. R. M. was supported by NSF CAREER awardAST 98-74670 and a Cottrell Scholarship from the ResearchCorporation.

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