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The Astronomical Journal, 123:346–361, 2002 January
# 2002. The American Astronomical Society. All rights reserved. Printed in U.S.A.
LARGE-SCALE EXTENDED EMISSION AROUND THE HELIX NEBULA: DUST, MOLECULES, ATOMS,
AND IONS
A. K. Speck, M. Meixner, D. Fong, P. R. McCullough,1 D. E. Moser, and T. Ueta
Department of Astronomy, University of Illinois at Urbana-Champaign, MC-221, 1002 West Green Street, Urbana, IL 61801
Received 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 (NGC
7293). The ionized gas is observed in the form of an H image, which is constructed using images from the
Southern H Sky Survey Atlas. The molecular emission was mapped using the H2 v = 1 ! 0 S(1) line at
2.122 lm. The far-infrared (FIR) observations were obtained using ISOPHOT on the Infrared Space Observatory. The H observations are more sensitive than previous measurements and show the huge extent of the
Helix, 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 previous 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 photodissociation
regions (PDRs) shows that the emission arises from thermal excitation of the hydrogen molecules in PDRs
and not from shocks. The FIR observations, at 90 and 160 lm, represent mostly contributions from thermal
dust emission from cool dust grains but include a small contribution from ionized atomic lines. Comparison
of the FIR emission with the H observation shows that the dust and ionized gas are coincident and extend
to 110000 radius. This equates to a spatial radial extent of more than 1 pc (assuming a distance to the Helix
of 200 pc). Assuming that the outer layers of the circumstellar shell have spherical symmetry, radiative
transfer modeling of the emission in H gives a shell mass of 1.5 M. However, the modeling does not cover
the 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 larger
than 3.5 lm. Comparison of these new observations with previous observations shows the large-scale stratification of the Helix in terms of ionized gas and dust, as well as the coexistence of molecular species inside the
ionized 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
The large-scale structure of the extended emission around
the Helix has been widely studied for the last fifty years. The
current paradigm for the structure of the nebula is that of a
disk composed of concentric rings in which the level of ionization decreases with distance from the central star. The
central zone, which appears to be a hole in the classic optical
image, is characterized by He ii emission, from recombinations of He2+ (O’Dell 1998), and [O iv] emission (Leene &
Pottasch 1987). Outside this central, highly ionized region
there is a ring of lower ionization species, that coincides
with the inner part of the optical ring and is characterized
by [O iii] and H emission (Warner & Rubin 1975; O’Dell
1998; Henry, Kwitter, & Dufour 1999). The H emission
persists beyond the [O iii] ring and also coincides with the
outer part of the optical ring characterized by the even lower
ionization species [O ii] and [N ii] (O’Dell 1998; Henry et al.
1999). To date most optical images do not show ionized
emission beyond the classic nebula. However, O’Dell (1998)
and Malin (1982) both show an arc of ionized emission out
at an angular radius of 85000 .
Observations of molecular species, CO (Huggins & Healy
1986, 1989; Healy & Huggins 1990; Forveille & Huggins
1991; 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 coincident with and does not extend beyond the ionized emission.
Molecules are expected to be destroyed by the passage of
the ionization front and excited in the photodissociation
1. INTRODUCTION
At a distance of only 200 50 pc (Harris et al. 1996;
Harrington & Dahn 1980), the Helix Nebula has a large
angular extent on the sky (140 radius at the largest
extent; O’Dell 1998; Malin 1982), making it extremely
valuable for the study of the spatial distributions of the
different species in nebulae (i.e., ions, atoms, molecules, and
dust). Understanding the distributions of the different
species enables a detailed analysis of the physical and chemical conditions during the postasymptotic giant branch
(post-AGB) phase and allows a better understanding of the
chemical enrichment of the interstellar medium (ISM) from
these intermediate-mass stellar sources.
The current work looks at the large-scale structure and
distribution of species across the nebula. We present new
observations at both near- and far-infrared (IR) wavelengths, showing the distribution of the molecular hydrogen
and cool dust grains, respectively, as well as a new deep H
image, which shows that the ionized part of the nebula has a
larger extent than previously seen. These new observations
are compared with those at other wavelengths from radio to
X-ray to obtain a comprehensive picture for the large-scale
structure of the Helix.
1
Cottrell Scholar of Research Corporation.
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EMISSION AROUND THE HELIX NEBULA
region (PDR) that exists at the boundary between the ionization front and the molecular gas. For this reason, it is
expected that there should be a ring of neutral and molecular species outside the rings of ionized gas. This is not seen
in the observations. The coincidence of molecular emission
with the ionized nebula implies that these species are somehow shielded from the ionizing Lyman continuum by being
inside dense (optically thick) condensations of material.
Indeed, evidence for such condensations has been known
for nearly fifty years (e.g., Zanstra 1955; Aller 1954), and
they are now known as the ‘‘ cometary knot’’ (e.g., O’Dell &
Handron 1996 and references therein). Dust exists throughout the nebula, as is shown in the IRAS far-IR images presented by Leene & Pottasch (1987).
In x 2 we present our new observations of the spatial
distribution of ionized gas (H), molecular hydrogen,
and cool dust. Section 3 compares the new observations
with the numerous studies at other wavelengths and dis-
347
cusses qualitatively the implications of the distributions
of the various species in terms of the large-scale structure of the nebula. In x 4 we show the results of modeling the nebula by using a simple spherically symmetric
radiative transfer model. The conclusions are discussed
in x 5.
2. OBSERVATIONS
We have observed the Helix in H ionized emission, farinfrared (FIR) cool dust continuum emission and near-IR
(NIR) molecular hydrogen lines.
2.1. H Observations
The continuum-subtracted, deconvolved H image of the
Helix, presented in Figure 1, was constructed from images
from 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 3 107 ergs cm2 s1 sr1. Higher contours correspond to
logarithmic 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.
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SPECK ET AL.
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, azimuthally
symmetric halo. This latter halo was removed by five iterations of a Lucy-Richardson deconvolution, which allowed a
better look at the point-symmetric structures extending out
from the Helix as far as 180 (see also Malin 1982).
The central ringlike structure in Figure 1 is the classic
optical ring of the Helix. The most remarkable aspects of
this new deep H image are the spatial extent of the emission and its point-symmetric structure. The ionized gas
associated with the Helix Nebula has an angular radius of
110000 from the central star. At a distance of 200 pc this
gives the Helix a physical radius of 1.1 pc. This image also
shows that the Helix is a density-bounded nebula (i.e., the
ionizing photons from the central star reach beyond the
edge of the material that makes up the nebula), with the nebular H emission merging with the background ISM H
emission. The shape of the extended ionized emission is
elongated in the east-west direction and shows evidence of
point-symmetric structures in the northeast-southwest,
east-west, and southeast-northwest directions, marked in
Figure 1 as A1–A2, B1–B2, and C1–C2, respectively. The
arc of ionized emission previously seen by O’Dell (1998) and
Malin (1982) can easily be seen to the northeast of the main
ring, within the extent of the point-symmetric structures.
To determine how much of the azimuthally symmetric
halo was due to the wings of the point-spread function
(PSF), we created an azimuthally symmetric PSF from hundreds of stars in scores of SHASSA images and deconvolved
the image using this PSF. The PSF of a very bright star such
as Sirius appears to be azimuthally symmetric (aside from
huge bleed trails), so we have assumed the PSF is azimuthally symmetric to get adequate signal for the PSF to
approximately r = 400 . The profile of the PSF is shown in
Figure 2, together with profiles of both the original and
deconvolved H images of the Helix.
Fig. 2.—Slice through the H image, together with the point-spread
function (PSF) of the H image, as determined from images of stars and
two PNs (NGC 7009 and NGC 6572). The two PNs have the same radial
profile as stars, even though the former are pure line emitters, whereas the
latter are continuum emitters. Thus, we can confidently use a PSF determined from many stars to deconvolve the line emission of the Helix. The
azimuthally symmetric PSF that we adopted is the dashed line; it has a
Gaussian core with FWHM = 1<3, an r2 region out to 110 , and an r3 profile beyond that. The Gaussian core-plus-r2 profile has been observed in
much larger Schmidt optics, albeit with different scale factors (King 1971).
The r3 outer profile avoids the logarithmic divergence of the integral to
infinity of an r2 profile.
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We prepared the image for the Lucy-Richardson deconvolution method by converting to photon units, forcing negative pixels (residuals around stars from continuum
subtraction) to zero, and adding in a DC level from telluric
emission lines. (The DC level was removed after the deconvolution). For the contour plot in Figure 1, we smoothed
the deconvolved image with a 3 pixel by 3 pixel median
filter.
2.2. ISOPHOT FIR Observations
The FIR observations were obtained using the Infrared
Space Observatory (ISO). We have obtained FIR linear
scans of the Helix Nebula using ISOPHOT (the imaging
spectro-photopolarimeter on ISO; Lemke et al. 1996).
These linear scans, which are centered on the central star of
the nebula, represent traverses across the entire circumstellar dust shells. The scans were taken on 1997 May 5 at a
position angle of 155 east of north. The position, orientation and extent of the linear scans relative to the optical
image of the Helix is shown in Figure 3. Both forward and
reverse scans along this track were obtained to determine
the repeatability of structures observed in the object’s spatial profile. The scans were imaged with two filters, the C100
90 lm, which has a 3 3 pixel format with 4600 pixel scale,
and the C200 160 lm, which has a 2 2 pixel format with
9200 pixel scale. The PHT32 AOT uses a combination of raster mapping and chopper sweeping to create a map. For
each raster step (of 6000 for the 90 lm and 9200 for the 160 lm
filter), the chopper was used to make smaller steps of 1500
and 3000 to give an image pixel size of 1500 4600 and
3000 9200 for the 90 and 160 lm filters, respectively. The linear scans used the maximum raster length possible of 300
and 460 for the 90 and 160 lm filters, respectively. The chopper, however, sweeps across the raster position, which
increases the total length of the linear scans to 360 for the 90
lm and 530 for the 160 lm filter. The point-spread functions
(PSFs) have full-width half-maxima (FWHM) of 44>5 for
the 90 lm filter and 97>4 for the 160 lm filter. The ISOPHOT data presented in this paper were reduced using the
PHOT Interactive Analysis package (Gabriel et al. 1997),
together with the P32Tools package developed specifically
for reduction of the AOT PHT32 data (Schulz & Peschke
2002). The P32Tools package allows a better calibration of
the PHT32 data, so that the errors on the calibrated data
can be reduced to 30%.
The flux-calibrated profiles of the ISO linear scans are
shown in Figure 4. We also show the ISO PHT22 data from
the ISO archive (principal investigator, Cox) for comparison. The PHT22 observations have a lower angular resolution and a smaller image area compared with our PHT32
observations. Slices through the PHT22 images at the same
position angle as the PHT32 linear scans are also included
in Figure 4. The errors on the photometry of the ISOPHOT
data are 30%. Therefore, the two 90 lm scans (PHT22 and
PHT32), shown in Figure 4 (top), are consistent within the
calibration errors. The PHT32 160 lm and PHT22 180 lm
scans show very similar morphologies. In fact the main difference between the 160 and 180 lm filters is that the 160 lm
filter is 20 lm broader to the short-wavelength end (see
Table 1), which accounts for the difference in brightness of
the nebula seen by these two filters. The 160 lm scan shows
evidence of low-level dust emission all the way out to an
angular radius of 110000 from the central star. This is also
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EMISSION AROUND THE HELIX NEBULA
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Fig. 4.—Profiles of the ISOPHOT PHT32 linear scans of the Helix,
together with slices through the lower resolution two-dimensional ISOPHOT PHT22 images taken at the same position angle as the linear scans.
The dotted lines represent the point-spread functions.
Fig. 3.—Extent, position and orientation of the ISO linear scans with
respect to the classic optical image of the Helix
seen in the IRAS data (see Fig. 5). Unfortunately the other
data (PHT22 90 and 180 lm, and PHT32 90 lm) do not
cover as large a distance from the central star, and therefore
the behavior of the outlying dust emission at these wavelengths is difficult to determine.
2.2.1. Comparison with the IRAS Images
Figure 5 shows the ISO PHT32 observations together
with the ISO PHT22 and IRAS 60 and 100 lm data. It is
clear that the ISO PHT22 and ISO PHT32 90 lm profiles
are very similar to the IRAS 60 and 100 lm profiles. The
morphology of the ISO PHT32 160 lm and the ISO PHT22
180 lm profiles, while similar to one another, show a
marked difference in comparison with the other data. They
are clearly missing the central region emission seen in the
shorter wavelength profiles (at 60, 90, and 100 lm). This is
discussed further in x 3.1.
The motivation for observing the Helix at the long wavelengths available to ISOPHOT was to determine the distribution of cool dust associated with the nebula. Before we
can properly compare the FIR images at different wavelengths we need to understand the contribution to the emission seen in these broadband images from the narrow
emission 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 following standard reduction procedures by using ISAP 2.0a.
The subspectra were merged to form a continuous spectrum
by comparing the continuum levels in the overlapping
regions. Line fluxes were measured by fitting a Gaussian to
the line profile.
We have PHT22 and PHT32 scans using three different
filters (90, 160, and 180 lm), all of which are potentially contaminated by line emission. Table 1 shows the emission lines
that fall within the filters of both the IRAS and the ISO
observations. The 90 lm filter covers the wavelength range
69–121 lm, which includes the [O iii] 88 lm fine-structure
line. The 160 lm filter covers the 129–219 lm range, and the
180 lm filter covers the 150–221 lm range, both of which
include the [C ii] 158 lm and the [N ii] 205 lm fine-structure
lines. Comparing the flux contained in the emission lines to
the flux in the continuum for these broadband filters, we
found 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 bands
we can obtain only rough estimates, since the LWS01 spectrum goes only to 196 lm, while these long-wavelength filters go out to 220 lm. However, we still find that the [C ii]
158 lm emission line contributes less than 20% to the
emission in these bands.
The broadband FIR images obtained by IRAS are also
potentially contaminated by nebular line emission. The 60
lm band covers the 30–84 lm wavelength range, and the
100 lm band covers the 70–140 lm range. The IRAS 100
lm image is very similar to that of the ISOPHOT 90 lm
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SPECK ET AL.
TABLE 1
Emission Lines within the ISO IRAS
Broad FIR Filters
Major Lines
Minor Lines
IRAS 60 lm, 30–84 lm Range
[S iii] 33 lm .........
[O iii] 52 lm.........
[S i] 56 lm ...........
[N iii] 57 lm ........
[O i] 63 lm...........
[Si ii] 35 lm
[Ne iii] 36 lm
[F iv] 44 lm
[Fe iii] 52 lm
[P ii] 61 lm
[F ii] 67 lm
[Si i] 68 lm
ISO 90 lm, 69–121 lm Range
[O iii] 88 lm.........
[Al i] 89 lm
[Fe iii] 105 lm
IRAS 100 lm, 70–140 lm Range
[O iii] 88 lm.........
[N ii] 122 lm........
[Al i] 89 lm
[Fe iii] 105 lm
[Si i] 130 lm
ISO 160 lm, 129–219 lm Range
[O i] 146 lm.........
[C ii] 158 lm ........
[N ii] 205 lm........
[Si i] 130 lm
ISO 180 lm, 150–221 lm Range
[C ii] 158 lm ........
[N ii] 205 lm........
images, as one would expect. The 60 lm image is not contaminated by the 88 lm [O iii] fine-structure line but is potentially contaminated by the [O iii] 52 lm and the [N iii] 57 lm
fine-structure emission lines, as well as the [S iii] 33 lm line.
We concur with Leene & Pottasch (1987) that the ionized
emission lines do not contribute significantly to the emission
in these broad FIR bands, with the majority of emission in
our observations coming from cool dust. The LWS spectra
were taken at the position of the optical ring of the Helix
and not of the central region; therefore these estimates hold
only for the ring region. The optical ring region of the Helix
is bright at all FIR wavelengths. However, the ISO LWS
spectra of the optical ring region do not show any signature
dust features or even a peak in the continuum suggesting a
range in dust grain sizes. Therefore, a particularly distinct
grain size population is needed to produce such a spectrum.
Further radiative transfer modeling is required to constrain
the grain size distribution.
2.3. Molecular Hydrogen Observations
To make the NIR H2 map of the Helix we used the NearInfrared Imager (NIRIM; Meixner, Young Owl, & Leach
1999) at the Mount Laguna 1 m telescope2 on 2000 August
2–5. The imager was used with the 200 pixel scale, giving the
256 256 pixel array an imaging area of 8<5 8<5. The
large angular size of the Helix Nebula demanded mosaic
2 Mount Laguna Observatory is jointly operated by San Diego State
University and University of Illinois at Urbana-Champaign.
Fig. 5.—Comparison of all FIR data. In each case for which a twodimensional image is available (i.e., the PHT22 data and the IRAS 60 and
100 lm images), the profile represents a slice through the image at the same
position as our ISOPHOT PHT32 linear scans. For each data set an offset
has been added for clarity: the IRAS 60 lm data have an offset of 36 MJy
sr1; IRAS 100 lm data have an offset of 12 MJy sr1; ISO PHT22 90 lm
data have an offset of 7 MJy sr1; ISO PHT32 90 lm data have an offset of
8 MJy sr1; and the ISO 160 and 180 lm data are not offset.
mapping. The mosaic was made from images taken at 13
positions in a diamond shape centered at the position of the
central star at (22h29m38 95, 20 500 13>5). Positions farther
from the central point were observed, but the entire molecular emission from the nebula was found to be contained
within the 13 position diamond mosaic whose maximum
extent was 12<25 from the central star position 2h29m3895,
20 500 13>5. Off positions were taken regularly at offsets of
0=5 from the central star. The mosaic map was imaged at
three wavelengths: the v = 1 ! 0 S(1) line of H2 at 2.122
lm, Br at 2.166 lm, and the broadband K0 at 2.12 lm.
Images taken with the K0 filter had a typical integration time
of 3 s, while the images taken using the H2 and Br filters
had typical integration times of 80 s. The data were reduced
and compiled using IRAF. Each exposure was flat-fielded
to eliminate large pixel-to-pixel sensitivity variations in the
detector array and was then sky-subtracted. Sky emission
maps were constructed by taking the average of several sky
frames obtained close to the target field both in time and
space. We used point sources in each field as reference
points to determine relative offsets between frames. Finally
the frames were co-added using the derived offsets. The
resulting H2 image is shown in Figure 6. The observations
were calibrated using standard stars: BS 8709, PLX 5546,
HD 162208, and HD 203856. These stars were observed
using the H2, K0 , and Br filters several times throughout
each night to account for changing sky conditions. The K-
Fig. 6.—Helix Nebula imaged in H2 2.122 lm v = 1 ! 0 S(1) line
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SPECK ET AL.
band fluxes for these standard stars were obtained from the
Elias standard-star list for HD 162208 and HD 203856
(Bouchet, Manfroid, & Schmider 1991 and Carter 1990 for
BS 8709; Veeder 1974 for PLX 5546). We did not correct for
the small changes in air mass because the standard stars BS
8709 and PLX 5546 are close to the Helix and the time
between consecutive observations was small.
The flux calibration gives a peak value for the H2 emission
of 3 104 ergs s1 cm2 sr1. The average brightness is
2 104 ergs s1 cm2 sr1. Our Br images show that
the Helix was undetected in this line, giving a 1 upper limit
to the Br emission brightness of 7 108 ergs s1 cm2
sr1.
The K0 images are consistent with all the nebular emission
in this filter coming from the H2 v = 1 ! 0 S(1) 2.122 lm
line. The K0 images also provide a 3 detection of the central star in this band. This shows that the flux of the central
star at K0 (2.12 lm) is 2 1013 ergs s1 cm2 lm1. The
central star of the Helix is believed to be a 123,000 K white
dwarf (Bohlin, Harrington, & Stecher 1982). Using the Vband magnitude for the Helix Nebula’s central star quoted
in Bohlin et al. (1982) to normalize a 123,000 K blackbody,
we found that the expected K0 flux density is 7 1013
ergs s1 cm2 lm1. Therefore the K0 flux we detect from the
central star is less than that expected from the 123,000 K
white dwarf and does not show evidence for an M dwarf
companion, as suggested by Guerrero et al. (2001) to
account for the hard X-ray source coincident with the
Helix’s central star.
3. STRUCTURE OF THE HELIX
Comparison of our new observations, together with previous work, shows that the Helix is composed of three distinct 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 highly
ionized gas (He ii and [O iv]), as well as some odd emission
at 60–100 lm FIR wavelengths, and appears to be a hole in
the 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 radius
of 9000 (e.g., O’Dell, Henney, & Burkert 2000). This zone is
completely filled with emission from He2+ (O’Dell 1998)
and O3+ (Leene & Pottasch 1987). Our observations show
that there is also H emission, albeit at a lower level than in
the ring (see also Henry et al. 1999; O’Dell 1998). The most
extraordinary observations of the central region are the FIR
data (IRAS and ISO). These show that there is significant
emission in the central zone at 60, 90, and 100 lm that is not
seen 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 contribution from the [C ii] 158 lm and [N ii] 205 lm emission
lines to the 160 and 180 lm scans, we can assume that the
emission in these bands arises solely from dust. Dust that
emits at 160 and 180 lm will also contribute to the continuum emission in the 60–100 lm bands. Let us also assume
that the same dust component is also the major contributor
to the flux in the IRAS 60 and 100 lm bands and the ISO 90
lm band. Therefore the excess emission in the central region
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He 0
He+
He2+
dust
❂
H+
O 3+
O 2+
O+
Molecular
Clumps
shadow
dust
H2
CO
H0 H
+
star
light
Fig. 7.—Schematic view of a density-bounded planetary nebula, where
both the ionization front and the shock front have passed through the entire
visible nebula. The edge of the nebula is where the density is so low that we
can no longer see the emission. The knots provide many mini-ionization
fronts throughout the nebula. The inset shows a magnified view of a single
globule with the ionization front and PDR around the edge. Note that dust
is present everywhere there is gas.
of the 90 lm scan, as seen in Figure 8, is not from the same
dust responsible for the emission coincident with the optical
ring at about 25000 .
To understand the difference between these two scans the
160 lm scan has been normalized to the 90 lm scan. The
normalized 160 lm scan is shown together with the 90 lm
scan in Figure 8 (top). In Figure 8 (bottom) the 160 lm scan
has been subtracted from the 90 lm scan, so that the excess
central emission at 90 lm is clear. Also shown are the IRAS
25 lm [O iv] data (Leene & Pottasch 1987) and the He ii
data from O’Dell (1998). This shows that the 90 lm excess is
spatially coincident with He2+ and O3+ ions and therefore
with the most ionized part of the nebula.
There are two possible explanations for this excess emission in the central zone: dust or FIR fine-structure emission
lines. First, let us consider dust. The IRAS 60 lm image has
the same morphology as the IRAS 100 lm image (Leene &
Pottasch 1987) and the ISO 90 lm data. Therefore, if the
excess emission is due to dust, it must also contribute a similar amount of emission in the 60 lm band as in the 90 and
100 lm bands. Furthermore, Cox et al. (1998) found that
the mid-IR (5–16 lm) emission from the Helix is almost
exclusively due to H2 rotational lines and argon and neon
emission lines, with no signature dust features and very little
underlying continuum. Therefore, if dust is responsible for
the excess emission, a grain population must exist that emits
strongly at both 60 and 90 lm, but not at either 15 or
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EMISSION AROUND THE HELIX NEBULA
Fig. 8.—FIR emission in the central region of the Helix: top, 90 lm scan
together with the normalized 160 lm scan; bottom, difference between the
90 and normalized 160 lm linear scans compared with slices through the
IRAS 25 lm data and the He ii data from O’Dell (1998) at the same position
angles as the ISOPHOT scan. The 25 lm and He ii data have been normalized to the ISO data. The He ii data have been convolved with a Gaussian
to approximate the resolution of the ISO 90 lm scan. The dips below zero
on either side of the central peak are a result of the different angular resolutions of the 90 and 160 lm observations.
160 lm. This narrow wavelength coverage for dust emission
seems peculiar. The nature of the implied dust grain size distribution is discussed further in x 4.
It is also possible that the excess central emission at 60
and 90 lm is caused by atomic or ionized emission lines.
The emission lines that fall into the broadband filters of
the ISO and IRAS images are listed in Table 1. The only
major atomic or ionized emission line that falls within
the wavelength range of the 90 lm filter is the [O iii] 88
lm line. This line also contaminates the IRAS 100 lm
images, while the other [O iii] fine-structure line at 52 lm
falls within the range of the 60 lm filter. Comparison
between our 90 lm scan and a slice through the [O iii]
5007 Å image from O’Dell (1998) taken at the same position angle as our ISO data shows that they have similar
bulk morphologies (Fig. 9). However, the inner region of
the 5007 Å image does not show the extra emission we
see at 90 lm (Fig. 8, bottom). The 5007 Å emission from
O2+ covers a larger angular extent than the 90 lm excess
emission and peaks at the inner edge of the classic optical
ring, while the 90 lm excess emission is strongest in the
central region of the Helix. Therefore, if the excess
central emission is to be attributed to [O iii] fine-structure
emission lines, there must be a mechanism by which the
spatial distribution from different lines of the same species can have markedly different morphologies. While it
353
Fig. 9.—Comparison of spatial distribution of emission from different
species: [O iii] 5007 Å is from O’Dell (1998); [N ii] 6584 Å and H are from
Henry et al. (1999); H2 emission profiles are from NIRIM (this paper) and
IRAS data; dust emission at 90 and 160 lm are from ISO (this paper).
may be possible to manipulate the relative intensities of
the optical and FIR fine-structure lines by using the
difference in the critical densities for these lines, the result
would be to increase the 5007 Å line relative to the 88
lm line, not vice versa. Therefore, it seems unlikely that
such huge differences in morphology as seen between the
5007 Å and 90 lm emission could be explained this way.
Therefore, we believe that the excess emission in the central zone is due primarily to dust. Model calculations of
dust emission are discussed further in x 4.
3.2. Ring
The ring region is characterized by the classic optical
image 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 from
different species. Ionized gas, dust, and molecular gas all
have their strongest emission in the ring region, although
the exact distributions vary slightly. The [O iii] 5007 Å emission is stronger in the central zone than the other ionized
species and drops off more quickly to the outer part of the
ring. The dust, as seen in the ISOPHOT linear scans, follows
the same distribution as the ionized gas (H). The dust and
H emission extends farther than that of the other species
shown in Figure 9. By comparing the new H2 observations
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Fig. 10.—Comparison of the extent of the molecular hydrogen with that of the ionized hydrogen
with new deep H observations (Fig. 10), it is clear that all
the molecular gas is entirely contained within the ionized
gas. This suggests that the molecular gas is shielded from
the ionizing radiation by being inside denser (optically
thick) condensations. Indeed such condensations have been
identified at the inner edge of the ring region and are known
as ‘‘ cometary knot’’ (see O’Dell & Handron 1996 and references therein).
3.2.1. Molecular Hydrogen Emission
H2 can be observed only where it is excited. H2 molecules
can be excited either collisionally, by shocks or high temperatures (thermally), or radiatively, through fluorescence.
Estimates of the H2 emission from PDR models led Cox et
al. (1998) to suggest that there must be some (shocked) collisionally excited H2 emission from the Helix to account for
its high intensity. However, PDRs in PNs are very different
from those associated with molecular clouds, to which most
published models refer. Most PDR models assume that an
equilibrium point is reached at which H2 formation is equal
to H2 destruction. However, for PNs, the PDRs are unlikely
to reach such an equilibrium since the timescale for H2 formation is generally long compared with the timescale for the
turn-on of the far-ultraviolet (FUV) flux (Sternberg 1998).
In this case the H2 emission may well be much higher than
expected from equilibrium PDR models (e.g., Sternberg
1998; Natta & Hollenbach 1998).
Natta & Hollenbach (1998) have modeled the emission
expected from PNs, including the evolution of the central
star and the ensuing PDR that propagates through the circumstellar shell. In their models the shock that propagates
through and excites the molecular shell is due to the acceleration of the superwind to 25 km s1 into the slower, less
dense AGB wind moving at 10 km s1. They do not
include the fast wind, since a shock produced by a fast wind
(velocity 1000 km s1) would dissociate the hydrogen molecules rather than excite them. Since there is currently no
evidence for a fast wind in the Helix (e.g., Patriarchi & Perinotto 1991) its possible effect will not be considered here.
Natta & Hollenbach (1998) also include the effect of X-rays
produced by hot central stars such as that of the Helix. Previous models generally ignore X-rays and include only the
FUV radiation. For a PN born of a high-mass progenitor,
X-rays make a significant contribution to the PDR emission, leading to much higher H2 emission than previously
expected. For their highest-mass progenitor (5 M) and
lowest-density nebula, the intensity of the H2 v = 1!0 S(1)
line is 105–104 ergs s1 cm2 sr1 after about 104 yr, consistent with but slightly lower than our observations, with a
peak of 3 104 ergs s1 cm2 sr1. The intensity is lower
for lower mass progenitors and higher density nebulae.
However, the progenitor mass for the Helix Nebula is
expected to be higher (6.5 M; e.g., Henry et al. 1999;
Gorny, Stasinska, & Tylenda 1997), which should yield a
higher flux from H2. Furthermore, while the models of
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EMISSION AROUND THE HELIX NEBULA
Natta & Hollenbach (1998) do not include clumpiness, they
suggest that further enhancement in the PDR H2 emission is
expected as a result of the photoevaporation of the cometary knots and filamentary condensations seen in the Helix.
This process effectively causes rapid advection of H2 from
within the condensations out into the mini-PDRs at the
knot 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 particularly for clumpy nebulae such as the Helix. Based on the
models of Natta & Hollenbach (1998) we argue that the high
level of H2 emission seen in the Helix is consistent with that
expected from the PDRs in an evolved PN, that consists of
the mini-PDRs that occur as a result of the clumpy nature
of the molecular shell.
3.2.2. Extent of the Molecular Emission
Since all the molecular emission appears to come from
cometary knots and related condensations, what does this
tell us about the evolution of this PN? Various formation
mechanisms for these knots have been put forward, starting
with Zanstra’s suggestion that they formed from clumps of
material already present in interstellar space. Since then the
formation mechanisms have fallen into two basic categories:
(1) clumps exist prior to the PN phase and the knots remain
after the passing of the ionization front and/or the fast wind
shock front (e.g., Zanstra 1955; Dyson et al. 1989) and (2)
the knots form after the onset of the PN phase, arising from
Rayleigh-Taylor (R-T) instabilities at either the ionization
front or the fast wind shock front (e.g., Mathews 1968; Capriotti 1971, 1973). At present there is no conclusive evidence
to support either mechanism outright. Whatever their formation mechanism may be, these dense condensations are
important as they are the reason we can still observe molecular emission from an evolved planetary nebula.
The inner edge of the optical nebula, at 10000 , where the
innermost cometary knots lie, marks how far the circumstellar shell that resulted from AGB mass loss has drifted since
the end of the AGB mass-loss phase. Assuming that the circumstellar shell simply expands with constant velocity away
from the star, this implies that the AGB dust shell has
expanded by 10000 in every direction since the end of the
AGB. The molecular emission, as seen in both H2 and CO,
currently extends to 40000 –50000 radius. Therefore, at the
end of the AGB the molecular shell would have extended
out to 30000 –40000 . At a distance of 200 pc, this equates to
a radial extent for the molecular envelope at the end of the
AGB of 0.3–0.4 pc (1018 cm).
According to Mamon, Glassgold, & Huggins (1988) the
extent of the molecular envelopes around cool, evolved stars
(such as AGB stars) is governed by photodissociation by the
interstellar radiation field (ISRF) and the ability of the circumstellar shell to shield itself from that radiation. They
found that, depending on the AGB mass-loss rate, the
extent of the CO envelope around AGB stars is expected to
be in the range 1016–1018cm. The upper limit is very similar
to the inferred extent of the molecular envelope of the Helix
at the end of the AGB.
However, the current molecular emission is easily seen to
originate from clumps that reside inside the ionized nebula,
which have allowed the survival of molecular species as the
ionization front passed through the nebula. If these clumps
are a result of the general AGB mass-loss processes, then
355
such clumps would also shield the molecular species from
dissociation by the ISRF, and the molecular envelope would
survive to a much larger radius than predicted by the simple
models of Mamon et al. (1988). Why, then, do we not see
molecular emission at a farther distance from the central
star?
There are three possible scenarios: (1) If the molecular
clumps are not self-gravitating they will gradually expand as
they drift away from the star and thus lose their ability to
shield themselves from the ionizing radiation. In this way
the clumps would be confined to a region relatively close to
the central star. (2) The molecular material does not form
clumps during the AGB phase and thus the extent of the
molecular shell is limited by the ISRF photodissociation as
modeled by Mamon et al. (1998). In this case, the clumpy
nature of the molecular emission as we now see it is due to
R-T instabilities caused by the passage of the shock front or
the ionization front through the relatively smooth molecular envelope. Since the smooth molecular envelope was limited to 1018 cm by the ISRF, this is also the extent of the
molecular material as it forms into clumps. However, this
assumes that the clump formation occurred more or less
instantaneously at the onset of the PN phase. (3) The
clumpy nature of the AGB mass loss manifests itself only
during the last stages of the AGB phase. Thus the earlier
molecular envelope was smooth and subject to limitation by
the ISRF, while the later clumpier molecular material is all
that has survived both the attack of the ISRF and the passing of the shock and ionization fronts with the onset of the
PN phase. The evidence for such clumpiness in the AGB
phase is seen in the maser spots, which do, indeed, occur in
the latest stages of AGB mass loss (see Dyson et al. 1989).
3.3. Outer Halo
The outer halo is characterized by low-level H and FIR
emission. Comparison of previous observations with our
new H and FIR data shows that the classic optical image is
merely the tip of the iceberg, with the ionized gas and dust
emission extending to a radius of 110000 (or 1.1 pc). The
relative distributions of the dust (as seen by the IRAS 100
lm data) and the ionized gas (as seen in H) are shown in
Figure 11, which shows the contours of IRAS 100 lm data
(processed through the HIRES algorithm to get better resolution; see Aumann, Fowler & Melnyk 1990) overlaid on
the H emission image. Figure 12 shows the profiles of the
emission in the ISO 160 lm and IRAS 60 lm and 100 lm
bands, together with the H emission profile. The profiles
have been normalized to the ISO 160 lm emission profile
for 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 of
north) is such that the outer, low-level dust emission seen in
Figure 12 coincides with the region of Figure 11 where the
dust emission appears to extend beyond the ionized gas
emission. However, the rapid drop-off with radius in ionized
gas emission relative to the dust emission is a result of the
emission mechanism and not an indicator that there is dust
without gas
R in these regions. Dust emission strength is given
by I / dl B(T), where dl is the pathlength, is
the density, is the emissivity of the dust grains, and
while the H emission
B(T) is the Planck function,
R
strength is given by IH / n2e dl, where ne is the electron
density. Therefore, H emission is more sensitive to the
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Vol. 123
Fig. 11.—Very extended 100 lm IRAS emission from the Helix compared with the H emission
decreasing density than the dust emission and drops off
more quickly. The dust emission is very sensitive to
temperature. At the temperatures relevant to this outer
halo region, small changes in dust grain temperature
can have a huge effect on the dust emission strength in the
FIR. The difference in the distributions of the H and the
FIR dust emission can be explained by their differing
dependence of density and temperature. It is also possible
that the Helix has a structure like that of some PNs and
proto–planetary nebulae (e.g., NGC 7027 and the Egg Nebula), with a bipolar structure superposed on spherical dust
shells (Hrivnak, Kwok, & Su 2001; Sahai et al. 1998; Latter
et al. 2000). The H extended emission shows three pointsymmetric pairs (marked in Figure 1 as A1–A2, B1–B2, and
C1–C2), reminiscent of these younger objects, that may be
due to precessing jets (e.g., Miranda, Guerrero & Torrelles
2001).
4. MODELING THE HELIX
4.1. Model
We have modeled the emission from various species
within the Helix using a simple radiative transfer model. We
refer the reader to Hoare (1990) for details concerning the
model. The radiative transfer assumes spherical symmetry,
a blackbody of Teff = 123,000 K (Bohlin et al. 1982), and a
constant gas-to-dust ratio throughout the nebula. The
model ionized gas emissions assume atomic abundances
from Henry et al. (1999). These abundances show that the
Helix has an approximately solar carbon-to-oxygen ratio
and is therefore oxygen-rich. For this reason the dust emission is assumed to arise predominantly from silicates. The
optical constants included in the model are for ‘‘ astronomical silicate ’’ from Draine & Lee (1984, 1987). The dust
grains are in thermal equilibrium with the radiation field
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EMISSION AROUND THE HELIX NEBULA
Fig. 12.—Very extended FIR emission from the Helix: profiles of the
ISO 160 lm (solid line), together with slices through the IRAS 60 and 100
lm bands and the H emission at the same position angles as the ISOPHOT scan. The profiles have been normalized to the ISO 160 lm emission
profile for better comparison.
357
that includes direct stellar radiation, free-free continuum
emissions, and line emissions (mostly Ly).
This simple model has several limitations. First, the
model assumes spherical symmetry, which is obviously a
problem when applied to an object such as the Helix, which
is patently not spherical. Second, the model assumes that
the gas and dust is smoothly distributed. Again, in the case
of the Helix, it is known that the nebula is not smooth, but
rather that it contains clumps of denser material (e.g., the
cometary knots). Finally, the largest grain radius that the
program can include is 3.5 lm. If larger grains are present in
the Helix their effect must be inferred. However, having laid
out the limitations and in the absence of something more
elaborate, using this simple model gives an indication of the
effect of the changing distributions of different ions, atoms,
and dust grains on the expected emission profiles. In particular it will help us to understand the very extended cool dust
emission and its variation with wavelength (see x 2.2.1) in
terms of dust grain size distributions and gas-to-dust ratios.
The model assumes that the dust and gas are coupled, so
that the density distribution of the dust follows that of the
gas. Therefore, the first step was to match the spatial distribution of the H emission.
4.2. Gas Density Distribution
Fig. 13.—Model emission profiles. Top, comparison of the model profiles of dust emission at 90 and 160 lm with the observed ISO linear scan
profiles, showing 90 lm observed data (asterisks) and 160 lm observed data
( plus signs). The 90 lm model profile has been normalized to the observation to show the shape of the distribution. Middle, observed profiles representing slices through the observed data taken at the same position angle as
the ISO linear scans. The solid line shows the best-fit model. Bottom, input
density profile.
We have two sets of H emission observations available.
Our own, new data (see x 2.1) is deeper and covers a larger
area of sky than that of Henry et al. (1999) but with lower
resolution. Therefore both are valuable in this study of the
spatial distribution of various species. Both data sets are
shown in Figure 13. The difference in the breadth of the
peak of emission is due to the difference in resolution, while
the difference in the extent of the observed emission is due to
the depth of the observations. Our data show that there is
low-level H emission extending to more than 30000 farther
than that seen in the data of Henry et al. (1999).3
The input profile that gives the best fit to the spatial distribution of H is shown in Figure 13, together with the
observed data. This profile is the result of 70 attempts at
matching the H profile, starting from a simple r2 density
drop-off. The best-fit profile has very low density in the central region, which is constant at 12 cm3 out to the edge of
the optical ring. From the edge of the optical ring (at 10000 )
to the peak of the density distribution at 30000 from the central star, the density rises steeply, as r2. The peak in the density profile is only 60 cm3. From the peak position the
drop-off in density is inversely proportional to the radius
out to a radius of 50000 , where the drop-off increases to r4.
The r4 dependence was found empirically by trying to balance having enough material to see the ionized emission but
not so much as to make the shell optically thick and thus
curtail the spatial extent of the H emission. This rapid
drop-off (faster than r2) is indicative of a superwind massloss episode in which the rapidly increasing mass loss leads
to such a distribution. It was not possible to find a density
distribution that would fit the outlying low-level H emission beyond 60000 . This is probably due to the spherical
geometry of the model.
The low level of the central density (12 cm3) is much
lower than expected following the work of O’Dell (1998)
3 In fact the low-level H emission extends as much as 60000 farther than
that seen by Henry et al. (1999), depending on direction; see Fig. 1.
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TABLE 2
Parameters for the Best-Fit Radiative
Transfer Model
Parameter
Value
T* .............................................
L* .............................................
D..............................................
Grain size distribution, ng.........
amin ..........................................
amax ..........................................
Peak density.............................
Gas-to-dust ratio .....................
Mass in circumstellar shell........
123,000 K
100 L
200 pc
a3.5
0.005 lm
3.5 lm
60 cm3
1000
1.3 M
and Henry et al. (1999) and leads to the He ii emission being
more extended than has been observed (O’Dell 1998). This
is probably due to the fact that our model assumes spherical
symmetry. If the Helix were more appropriately modeled as
a disk, it would be possible to include more material in the
central zone. This would produce a nebula that is ionization
bounded in the plane of the disk but density bounded perpendicular to the disk, giving rise to the confined He ii emission seen in the disk but allowing for the very extended H
emission. This suggests that the extended H emission is a
projection of a bipolar flow, with the optical ring as the
waist (see O’Dell 1998).
The Strömgren radius for hydrogen for our best-fit model
is beyond the edge of the shell, producing a density-bounded
nebula. Table 2 lists the important parameters for the bestfit model.
4.3. Modeling the Dust
Having fitted the shape of the H emission reasonably
well, we concentrated on the effect of the grain size distribution and gas-to-dust ratio on the emission at FIR wavelengths where the majority of the emission comes from cool
dust. These models are compared with the ISO observed
data.
For each set of model input parameters the grain size distribution is described by ng = a3.5, with the minimum
(amin) and maximum (amax) grain sizes being specified for
each model. Figure 14 shows the effects of changing the
grain size limits. When only small grains (0.005–0.05 lm)
are used, the emission very close to the central star is too
high at all IR wavelengths. While high central emission is
expected at 60 and 90 lm, this central emission is not
observed for mid-IR wavelengths (e.g., Cox et al. 1998) or
at 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 the
observed data in the central region. While there is still a little
more central emission in the model profiles than observed,
this is partly due to the projection of the spherical shape of
the model onto the plane of the sky. However, in the case of
large grains, the model profile falls far short of fitting the
outer wings of the FIR emission. Unfortunately the model
is not capable of including different grain size distributions
in different regions. Changing the grain size distributions
implies that there is a paucity of small grains in the central
regions for the dust emission strength to remain low in the
mid-IR but that small grains are needed in the outer parts of
the optical ring and beyond to produce the high level of FIR
Fig. 14.—Effect of the grain size distribution on the spatial profile of the
infrared emission. Top, 90 lm profiles. The asterisks show the profile of the
ISO 90 lm linear scan. Bottom, 160 lm profiles. The asterisks show the profile of the ISO 160 lm linear scan. The model profiles are normalized to the
ring (25000 ) in the observed emission profiles.
emission seen beyond 30000 . The very low level emission
observed in the ISO 160 lm and IRAS 100 lm data beyond
80000 is probably partly heated by the ISRF (see YPK; the
Egg Nebula and CRL 618, Speck, Meixner, & Knapp 2000).
It is also possible that the clumpy nature of the Helix plays a
role here. The model assumes that the shell is smoothly distributed, but the clumpy nature of the Helix may allow photons to escape through the gaps between the dense clumps
and thus the ionized emission would not be as confined as
the model suggests. This could also contribute to the heating
of dust grains in the outer halo.
For the best-fit model the widest dust grain size distribution was used (0.005–3.5 lm). However, the upper limit for
the grain size is set by the capabilities of the model. A much
larger maximum grain size is necessary to account for the
absolute flux-calibrated brightnesses in the observed FIR
data. By increasing the maximum grain size and maintaining the gas-to-dust ratio (by mass), the number of smaller
grains is decreased, and the mass from these grains goes into
the larger grains. Decreasing the number of small grains
lowers the overall flux levels in the IR. To match the
observed brightnesses at 90 and 160 lm for the Helix by
using the 0.005–3.5 lm grain size distribution, the gas-todust ratio must be increased to greater than 1000 (much
higher than the 100–200 usually assumed; see Fig. 15). The
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EMISSION AROUND THE HELIX NEBULA
Fig. 15.—Effect of the gas-to-dust ratio on the absolute brightness of the
infrared emission. The asterisks show the flux-calibrated profile of the ISO
160 lm linear scan.
absolute value of the brightness is also decreased by removing the small grains, but, as explained above, some small
grains are necessary to match the spatial distribution of the
FIR emission. Therefore, including very large grains (>3.5
lm) has the effect of decreasing the number of small grains
without changing the gas-to-dust ratio or the shape of the
emission profile at these wavelengths. To determine whether
there are very large dust grains in the Helix, observations
must be made at even longer wavelengths (submillimeter
and millimeter). Indeed, centimeter-sized grains have been
seen in proto–planetary nebulae (e.g., the Egg Nebula; Jura
et al. 2000). If large grains are not present, this implies that
the gas-to-dust ratio is very high and that there is much less
dust in the Helix than previously thought.
One odd aspect of the models is the behavior of the emission within the central zone. One of our objectives with
these models was to understand the origin of the central
emission seen at 60 and 90 lm and not at longer (160 and
180 lm) or much shorter (<25 lm) wavelengths. Figure 16
shows that in the central zone the dust heating is dominated
359
by the central star, while the dust heating in the optical ring
is dominated by Ly photons. Therefore, the extra emission
in the central zone comes from heating of dust by the central
star. However, the inclusion of material in the central region
manifests itself as a point source centered on the central star.
No dust density distribution could be found that would
cause the excess central emission to appear as more than a
simple central point source. As explained above, the central
zone needs to have a deficit of small grains relative to the
outer parts of the optical ring. Careful manipulation of the
grain size distribution allows the appearance of the (point
source) excess emission at 60 and 90 lm over and above that
seen at 25 and 160 lm. By restricting the grain sizes in the
central zone to the large grains (>1 lm), the central zone
may show an observed excess at 60 and 90 lm without
showing extra emission at 25 and 160 lm. However, to
entirely fill the central zone with such dust emission a
particularly distinctive size distribution that changes with
radius must be invoked. Very close to the star, there can be
only very few small grains. The relative number of small
grains needs to increase with distance from the star, out to
the edge of the optical ring. The increasing number of small
grains may allow the excess emission to fill the central zone,
rather than being confined to a point source. At the edge
of the optical ring the line heating, dominated by Ly,
begins to dominate the dust heating so that the dust
temperature increases again and the peak of the dust
emission is coincident with the peak of the H emission.
What mechanism can be responsible for the strange dust
grain size distribution necessary to produce the excess 60
and 90 lm emission? Since the relative number of small
grains needs to increase with distance from the central stars
(at least as far as the edge of the optical ring), we propose
the following mechanism: The high-energy photons produced by the hot central star have sufficient energy to
destroy small grains. With increasing distance from the star
the destruction of small grains by high-energy photons
decreases, producing the necessary grain size distribution.
4.4. Mass of the Nebula
Based on the model, the circumstellar shell mass for the
Helix is 1.3 M. The model requires a gas-to-dust ratio of
1000, but even with a more usual gas-to-dust ratio of 100,
the contribution to the nebula mass from the dust is negligible. The mass of the molecular gas is 0.025 M, and the
mass of neutral gas is 0.18 M (Young et al. 1999), giving a
total nebular mass of 1.5 M. However, the model
accounts for only the inner 50000 –60000 radial extent. The
low-level emission out to more than 110000 will contribute
significantly to the shell mass, making this a lower limit.
5. CONCLUSIONS
Fig. 16.—Relative contributions of the direct stellar radiation and the
H to the dust heating. These relative contributions are virtually identical
for all grain sizes within the distribution. The contribution of the diffuse
scattered starlight is always negligible.
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 than
previous measurements, revealing the huge extent of the
Helix Nebula, as well as confirming it as (at least partially)
density bounded. The H emission extends to a maximum
angular radial extent of 110000 , where it merges with the
background ISM H emission. At a distance of 200 pc,
this equates to a spatial extent for the ionized gas of 1.1 pc.
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SPECK ET AL.
Furthermore, there is evidence of point-symmetric structures in the extended H emission.
The ISOPHOT FIR observations have contributions
from both ionized atomic lines and thermal emission from
cool dust grains. We have shown that the contribution to
the 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 suggests that there is a population of relatively large dust grains
and a paucity of small grains within the central region of the
Helix. Moreover, to fit the shape of the observed emission,
the grain size distribution must include an increasing proportion of smaller grains with increasing distance from the
star. This implies that high-energy photons from the hot
central star have destroyed small dust grains close to the
star. Small grains are needed in the outer parts of the dust
shell to account for the shape and extent of the FIR emission at larger distances from the star. To match the flux levels of the FIR emission, our models needed gas-to-dust
ratios of 1000. This is much larger than the value 100–
200 usually assumed. One possible explanation is that much
of the dust mass is hidden in very large dust grains (>3.5
lm). Therefore, our results imply either that 80%–90% of
the mass of the dust is contained in dust grains larger than
3.5 lm in radius or that there is much less dust in the Helix
than previously thought. The cool dust emission has been
shown to extend to a radius of 110000 . At the distance of
the Helix this equates to a spatial radial extent of over 1 pc.
Our models of the Helix account for emission out to only a
radius of 50000 –60000 . There are two possible explanations
for the observed outer emission, which is not seen in the
models. At such a great distance from the central star, the
outer part of the dust shell may be partly heated by the
ISRF (see Speck et al. 2000; YPK). This additional heating
source is not included in our model. Alternatively, the
model assumes that the shell is smoothly distributed (i.e.,
not clumpy), whereas the Helix is obviously clumpy. It is
possible that the photons currently confined within the classic ring region may be able to reach greater radii by escaping
through the gaps between clumps. This would provide some
extra heating to account for the more extended emission.
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The H2 observations show that the molecular gas is
entirely contained within the ring region of the Helix, coincident with the bright ionized gas emission. Therefore all the
molecular gas is contained within clumps and cometary
knots in the ring region. The origin of these clumps is still
not clear. Comparison of the H2 emission strength with
time-dependent models for PDRs (Natta & Hollenbach
1998) shows that the excitation of the H2 molecules is
consistent with such models, especially if the clumpiness of
the molecular medium is included. It is not necessary to
invoke shocks to achieve the strong molecular emission
observed.
Based on a spherically symmetric model for the Helix we
get a circumstellar shell mass of 1.5 M. However, the
model accounts for only the inner 50000 –60000 radial extent.
The low-level emission out to greater than 110000 will contribute significantly to the shell mass, making this a lower
limit.
We are very grateful to Bob O’Dell, both for the use of his
data (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 ISOPHOT team (Carlos Gabriel, Rene Laureijs, Sybille
Peschke, and Bernard Schulz) at Vilspa, Spain, for their
help and patience with understanding the data reduction
and calibration techniques. We are also grateful to Mike
Barlow for numerous constructive conversations or
correspondences, and to Elric Whittington for reading
through the manuscript. The H data are from the
Southern H-Alpha Sky Survey Atlas, which was produced
with support from the National Science Foundation.
A. K. S. was supported by NASA JPL 961504 and NASA
STI 7898.02-96A. M. M., T. U., and D. E. M. were
supported by NSF CAREER award AST 97-33697.
D. F. was supported by The Laboratory for Astronomical
Imaging at the University of Illinois and NSF grant 9981363. P. R. M. was supported by NSF CAREER award
AST 98-74670 and a Cottrell Scholarship from the Research
Corporation.
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