The chemical composition of TS 01, the most oxygen-deficient planetary nebula

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A&A 511, A44 (2010)
DOI: 10.1051/0004-6361/200912405
c ? ESO 2010
Astronomy
&
Astrophysics
The chemical composition of TS01, the most oxygen-deficient
planetary nebula
AGB nucleosynthesis in a metal-poor binary star?,??,???
G. Stasi´ nska1, C. Morisset2, G. Tovmassian3, T. Rauch4, M. G. Richer3, M. Peña2, R. Szczerba5, T. Decressin6,
C. Charbonnel7, L. Yungelson8, R. Napiwotzki9, S. Simón-Díaz10, and L. Jamet1
1LUTH, Observatoire de Paris, CNRS, Université Paris Diderot, Place Jules Janssen, 92190 Meudon, France
e-mail: grazyna.stasinska@obspm.fr
2Instituto de Astronomia, Universidad Nacional Autonoma de Mexico, Apdo. Postal 70264, Mexico D.F., 04510 Mexico
3Instituto de Astronomia, Universidad Nacional Autonoma de Mexico, Apdo. Postal 877, Ensenada, Baja California, 22800 Mexico
4Institute for Astronomy and Astrophysics, Kepler Center for Astro and Particle Physics, Eberhard Karls University, Sand 1,
72076 Tübingen, Germany
5N. Copernicus Astronomical Center, Rabia´ nska 8, 87-100 Toru´ n, Poland
6Argelander Institute for Astronomy (AIfA), Auf dem Hügel 71, 53121 Bonn, Germany
7Geneva Observatory, University of Geneva, ch. des Maillettes 51, 1290 Sauverny, Switzerland and Laboratoire d’Astrophysique
de Toulouse-Tarbes, CNRS UMR 5572, Université de Toulouse, 14, Av. E.Belin, 31400 Toulouse, France
8Institute of Astronomy of the Russian Academy of Sciences, 48 Pyatniskaya Str., 119017 Moscow, Russia
9Centre for Astrophysics Research, University of Hertfordshire, College Lane, Hatfield AL109AB, UK
10Instituto de Astrofísica de Canarias, 38200 La Laguna, Tenerife, Spain
Received 30 April 2009 / Accepted 15 December 2009
ABSTRACT
The planetary nebula TS01 (also called PNG135.9+55.9 or SBS1150+599A) with its record-holding low oxygen abundance and its
double degenerate close binary core (period 3.9 h) is an exceptional object located in the Galactic halo.
We have secured observational data in a complete wavelength range to pin down the abundances of half a dozen elements in the
nebula. The abundances are obtained via detailed photoionization modelling which takes into account all the observational constraints
(including geometry and aperture effects) using the pseudo-3D photoionization code Cloudy_3D. The spectral energy distribution of
the ionizing radiation is taken from appropriate model atmospheres. Incidentally we find from the new observational constraints that
both stellar components contribute to the ionization: the “cool” one provides the bulk of hydrogen ionization, while the “hot” one is
responsible for the presence of the most highly charged ions, which explains why previous attempts to model the nebula experienced
difficulties.
Thenebular abundances ofC,N,O,andNearefoundtobe1/3.5, 1/4.2, 1/70, and1/11of thesolar valuerespectively, withuncertainties
of a factor 2. Thus the extreme O deficiency of this object is confirmed. The abundances of S and Ar are less than 1/30 of solar. The
abundance of He relative to H is 0.089 ± 0.009.
Standard models of stellar evolution and nucleosynthesis cannot explain the abundance pattern observed in the nebula. To obtain an
extreme oxygen deficiency in a star whose progenitor has an initial mass of about 1 M?requires an additional mixing process, which
can be induced by stellar rotation and/or by the presence of the close companion. We have computed a stellar model with an initial
mass of 1 M?, appropriate metallicity, and initial rotation of 100 km s−1, and find that rotation greatly improves the agreement between
the predicted and observed abundances.
Key words. planetary nebulae: individual: TS 01 – ISM: abundances – stars: AGB and post-AGB – binaries: general –
nuclear reactions, nucleosynthesis, abundances
?Based on observations obtained at the Canada-France-Hawaii
Telescope (CFHT) which is operated by the National Research Council
of Canada, the Institut National des Sciences de l’Univers of the Centre
National de la Recherche Scientifique of France, and the University of
Hawaii.
??Based on observations with the NASA/ESA Hubble Space
Telescope, obtained at the Space Telescope Science Institute, which is
operated by the Association of Universities for Research in Astronomy,
Inc., under NASA contract NAS 5-26555.
???Based on observations made with the Spitzer Space Telescope,
which is operated by the Jet Propulsion Laboratory, California Institute
of Technology, under NASA contract 1407.
1. Introduction
SBS1150+599A was discovered in the second Byurakan Sky
Survey and first classified as a cataclysmic variable (Stepanian
et al. 1999). Tovmassian et al. (2001) discussed in detail the
nature of the object and arrived at the conclusion that it is
in fact a planetary nebula (PN). The object was renamed
PNG135.9+55.9, following the nomenclature for Galactic
PNe from the Strasbourg-ESO catalogue of Galactic Planetary
Nebulae (Acker et al. 1992). For the sake of brevity we will re-
fer to it as TS01 in the rest of the paper. This PN is special in at
least three important aspects. First of all, its oxygen abundance
is very low, significantly lower than in any other PN known up
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A&A 511, A44 (2010)
to now (Tovmassian et al. 2001; Richer et al. 2002; Jacoby et al.
2002; Péquignot & Tsamis 2005). Second, its nucleus is a spec-
troscopic binary with a period of only a few hours (Tovmassian
et al. 2004). Third, it appears from estimates of the nature and
masses of the two stellar components, that TS01 could turn into
a double degenerate type Ia supernova (Tovmassian et al. 2004).
Each of these aspects, even taken alone, makes TS01 an excep-
tional object.
In this paper, we reexamine the chemical composition of
TS01. Briefly, the story of the determination of the chemical
composition of this object is the following. Tovmassian et al.
(2001) had optical spectra of TS01 in the range 3900−7000Å
obtained with 2m class telescopes which showed no lines from
heavy elements except a very weak [O iii] λ5007, with an in-
tensity a few percent of Hβ. A coarse photoionization analysis
suggested an oxygen abundance smaller than 1/100 solar. Note
that standard empirical methods for abundance determinations
in PNe cannot be used for TS01, since the electron temperature
cannot be determined directly from observations. To go further
in the abundance determination of TS01 required an estimate of
the effective temperature of the central star. One way is to ob-
tain a good blue spectrum of the PN, and use the [Ne v] λ3426/
[Ne iii] λ3869 ratio (or a limit on it) as a constraint. Richer
et al. (2002) at the Canada-France-Hawaii Telescope (CFHT)
andJacobyetal.(2002)attheMultipleMirrorTelescope(MMT)
secured deep blue spectra to detect these lines. Jacoby et al.
(2002)detectedthe[Nev]λ3426lineat alevelof0.8Hβ. Richer
et al. (2002) found only an upper limit of 0.1 Hβ! Concerning
the [Ne iii] λ3869 line, Jacoby et al. (2002) measured an inten-
sity about ten times higher than Richer et al. (2002). The two
papers appeared within a few days of each other on astro-ph,
revealing this big conflict in the observations. The two groups
conducted independent photoionization analyses, and both con-
cluded that the O/H ratio is less than 1/100 solar (the main rea-
son for their similar result for the oxygen abundance was the
similar [Ne v] λ3426/[Ne iii] λ3869 ratio used by both stud-
ies). Péquignot & Tsamis (2005) made a combination of the two
observational data sets and conducted their own photoioniza-
tion analysis. They concluded that the O/H ratio of TS01 lies
between 1/30−1/15 solar (still holding the record for the most
oxygen poor planetary nebula, but much higher than previously
published). However, Péquignot & Tsamis (2005) neglected to
consider observations of TS01 made with the Hubble Space
Telescope (HST) and the Far Ultraviolet Spectroscopic Explorer
(FUSE).As a result,someoftheir“predicted”lineintensitiesare
in conflict with what is actually observed in the UV. HST obser-
vations were obtained in 2003 and presented in a short, prelim-
inary version by Jacoby et al. (2006). Those authors quoted an
oxygen abundance of 1/30−1/40 solar, and carbon and nitrogen
abundances roughly 1/10 solar.
Before embarking on a new determination of abundances,
we have chosen to gather the best possible observations at all
wavelength ranges. These data provide many more constraints
than were available in any previous study. In order to make the
best use of the large amount of data obtained with different tele-
scopes, we use a pseudo-3D photoionization code, Cloudy_3D,
which is able to account for the nebular geometry as we see
it now, and with which we can properly take into account the
aperture effects. This code is based on CLOUDY (Ferland et al.
1998) and was written by Morisset (2006).
Thepaperis organizedas follows.Section2 presentsthe new
observationalmaterial:severalopticalspectra, HST imagingand
spectroscopy, infrared spectroscopy with the Spitzer Telescope,
and mentions our X-ray observations with XMM. Section 3
Table 1. Log and characteristics of the spectroscopic observations.
Telescope
FUSE
HST STIS
CFHT MOS
CFHT MOS
Kitt Peak
SDSS
Spitzer IRS SH 22 Apr. 06 9.9–19.6μm
Spitzer IRS LH 22 Apr. 06 18.7–37.2μm 600
Date
30 Jan. 02 900–1200 Å
4 May 03 1170–1700 Å 1.20Å
1 May 03 3400–5300 Å 3.0–3.5 Å
4 Mar. 01 3400–8000 Å 23 Å
1 Jan. 033600–7500 Å 7Å
3819–9196 Å 2–4 Å
λ RangeResolution Aperture
30 × 30??
0.5??
1??
5??
2??
diameter 3??
4.7 × 11.3??
11.1 × 22.3??
600
summarizes other data that we used as constraints for the pho-
toionization modelling. Section 4 describes our modelling strat-
egy and presents our “reference model”. Section 5 evaluates the
error bars on the derived elemental abundances, taking into ac-
count observational uncertainties in emission-line fluxes, uncer-
tainties in model input parameters and also uncertainties arising
from an imperfectdescriptionof the physical processes included
in the models. In Sect. 6, we compare the chemical composition
ofTS01withthatofotherPNe intheGalactichaloanddiscuss it
in terms ofstellar nucleosynthesisin theasymptoticgiantbranch
(AGB) phase. Finally, Sect. 7 summarizes our main findings.
2. New observational data on emission lines
We present the observational data that we secured on TS01 and
its stellar core since the work presented in Tovmassian et al.
(2004). Some of those data were already briefly reported in con-
ference proceedings, but here we describe the acquisition and
reduction processes in more detail. Note that the observations
and reductions were done by different people and at different
epochs, when our knowledge on the object was not the same.
This explains the differences in the tactics employed to reduce
the data, estimate the line fluxes and correct for reddening. We
did not try to fully homogenize the data reduction process, since
we felt it unnecessary for our purposes.
The log and characteristics of each set of observations are
given in Table 1.
2.1. Imaging
We (M.P.) retrieved the data corresponding to the proposal
ID 9466 from the HST archives and analyzed them. The obser-
vations were performed on May 5, 2003. Two types of data are
available: direct imaging and spectroscopy.
Direct imaging was obtained with the Advanced Camera
for Surveys (ACS). The high resolution channel with a field of
view of 26??× 29??and a plate-scale of 0.027??per pixel, with
filters around Hα (central wavelength 6581.97 ± 162.8Å) and
[Ne v] λ3426 (central wavelength 3432.8± 42.66Å) were used.
Figure 1 shows the Hα image obtained by averagingthe four
calibratedframes j8do01021,j8do01022,j8do01023,j8do01024
(870 s exposure time each; 58 min in total), after aligning them
with respect to j8do01021. The image is roughly elliptical in
shape with two brighter, symmetrically-placed lobes at a posi-
tion angle of about 103◦that extend the full major axis of the
ellipse. The nebula is not perfectly symmetric, with the outer-
most southern part much fainter. The size of the nebular image
is about 5??.
Figure 2 shows the same image as Fig. 1, with the differ-
ent observingaperturesindicated:continuouslinesforKitt Peak,
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G. Stasi´ nska et al.: The chemical composition of TS01
25.2s24.9s11h53m24.7s24.4s 24.2
59d40m01.0s
59.2s
57.4s
55.6s
Fig.1. HST-ACS Hα image of TS01, with a logarithmic grey scale. A
grid with the position scale (RA and Dec) is traced. North is up.
50100150
- - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - -
- - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - -
Fig.2. HST-ACS Hα image of TS01. The locations of the different
spectroscopic apertures are indicated: Kitt Peak (continuous lines),
HST STIS (dotted lines), SDSS (circle). See text for the CFHT 2003
observations.
dotted lines for HST STIS, circle for SDSS. For the CFHT 2003
observations, the slit was rotated before each of the seven expo-
sures (see Sect. 2.2.1), so as to remain as close as possible to the
parallactic angle. The position of the slit is not indicated in the
figure forthe sake ofclarity,but was taken into accountcorrectly
when comparing the predicted line intensities with the observed
ones (see Sect. 4.2).
Figure3 shows an averageof the [Nev]λ3432calibratedim-
ages j8do01011, j8do01012, j8do01013 and j8do01014 (3000 s
of exposure time each; 200 min in total). This image reveals a
very faint, roughly spherical extended nebulosity and an impor-
tant emission in the centre, probably caused by the stellar core.
Some faint extensions are marginally detected in the directions
of the Hα lobes.
2.2. Optical spectroscopy
2.2.1. CFHT data
CFHT 2003 TS01 was reobserved at the CFHT by M.R. and
G.S. on 1 May 2003 using the MOS spectrograph and a 1??
slit (Le Fèvre et al. 1994). The U900 grism was used, giving a
Fig.3. HST-ACS image of TS01 in the [Ne v] λ3426 line, on a loga-
rithmic grey scale. Same orientation and scale as Fig. 1.
spectral range of 3400−5300Å and a spectral resolution of
3−3.5Å (measured from arc lamp spectra). Seven 1800s expo-
sures were obtained. During each exposure, the slit was set to
within 10◦of the parallactic angle. Details of the reduction pro-
cess of the individual exposures are given in Tovmassian et al.
(2004).
To obtain a high signal-to-noise spectrum of the nebular
emission lines, it is necessary to account for the stellar and neb-
ular continuum emission. These contributions were subtracted
from the individual exposures before summing the individual
spectra. First, the observed spectra were shifted in velocity so
that the stellar absorption line was at rest. Next, the Hβ intensity
was measured and used to scale a model of the nebular contin-
uumemission.SincetheHβ fluxis affectedbystellar absorption,
the stellar absorption was assumed to have an equivalent width
of 13Å, a value typical for the models used (see below). This
scaled nebularcontinuumwas thensubtractedfromthe observed
spectrum. Then, a model stellar atmosphere was scaled so as
to match the observed continuum level and subtracted from the
observed spectrum. This procedure leaves a pure nebular emis-
sion line spectrum, supposing that the model nebular and stellar
continua are representative of their real counterparts. It is un-
likely that subtracting the continua introduces significant uncer-
tainty into our final line intensities. Model stellar spectra with
(Teff,logg) pairs of (90kK, 5.05), (120kK, 5.35), and (150kK,
5.56) computed by R.N. (see Tovmassian et al. 2004) were sub-
tracted from our observed spectra and the differences in the re-
sultinglineintensities werealways smallerthantheuncertainties
in the fits1.
Once the stellar and nebular continua were subtracted, the
individual nebular spectra were shifted back to their original ve-
locities and summed. We measured the nebular emission line
strengthsfromthis final spectrum.Thelineintensities weremea-
sured using INTENS, a locally-implemented software package
(McCall et al. 1985). This software simultaneously fits a sam-
pled Gaussian function to the emission line(s) and a straight
1This treatment was applied well before we had understood that the
optical continuum was dominated by a star of 55kK (see Sect. 3.3). In
view of the fact that the lines we use for the dignostics discussed in the
present paper are hardly affected by this correction, we decided not to
redo the subtraction using more adequate model stellar spectra.
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A&A 511, A44 (2010)
Fig.4. Flux-calibrated CFHT spectra of TS01. In blue is the CFHT
2003 spectrum with the nebular and stellar continua subtracted, in
black the original CFHT 2001 spectrum, and in red the CFHT 2001
spectrum with the correct wavelength solution.
line to the continuum. It returns the line strengths, line wave-
lengths, and uncertainties in these quantities. The line intensities
presented in Table 2 together with their uncertainties are those
measuredafter subtractingthe stellar spectrumfor Teff= 120kK
and logg = 5.35. The listed intensities are not corrected for red-
dening. In the case of no detection, two-sigma upper limits are
given instead.
We note that the [Ne v] λ3426 line is present, and strong.
Its intensity is of the same order as in the spectrum of Jacoby
et al. (2002), and much higher than the upper limit given by
Richeret al.(2002).Theremaininglineshaveintensities roughly
in agreement with those published by Richer et al. (2002) and
Jacoby et al. (2002), except for the [Ne iii] line which appeared
on the top of a bump in Jacoby et al. (2002) and was attributed a
high intensity in that paper.
CFHT 2001 In view of the important discrepancy with the
Richer et al. (2002) data concerning the [Ne v] λ3426 line, we
decided to reanalyze the spectrum of TS01 we had obtained in
March 2001 at the CFHT. First, though, we refer the interested
reader to Richer et al. (2002) for a discussion of the details of
these observations. The basic difficulty with these observations
was that the arc lamp spectra were taken with the same 5??slit
used for the object observations, which resulted in an arc spec-
trum with severelyblendedlines. Afterrepeatingthe wavelength
calibration more carefully, we found that our previous solution
had stretched out the spectrum at the shortest wavelengths, lead-
ing us to not recognize the [Ne v] λ3426 line because of its low
contrast with respect to the continuum and its erroneous wave-
length. The analysis of this spectrum was considerably simpler
given that the stellar features were not resolved. We simply fit
the continuum shape and removed it (INTENS assumes that the
continuum is a straight line), then measured the line intensities
with INTENS. The resulting line intensities are given in Table 2.
The intensity we find now for [Nev] λ3426 agrees with the one
given by Jacoby et al. (2002) and is compatible within two sig-
mas with the one obtained with the CFHT 2003 data. Slit effects
could perhaps explain the slight difference in intensity between
CFHT 2001 and CFHT 2003.
Table 2. Intensities of optical lines, corrected for stellar absorption, but
not for reddening, with respect to Hβ = 100.
SDSSKitt PeakCFHT 2003
55.54 ± 5.19 84.05 ± 13.07
<0.60
1.07 ± 0.31
1.70 ± 0.32
2.30 ± 0.35
2.64 ± 0.38
4.67 ± 0.37
0.77 ± 0.31
9.28 ± 0.45
12.70 ± 0.47
0.56 ± 0.21
<0.58
25.16 ± 0.43
1.04 ± 0.19
<0.32
47.19 ± 0.50
<0.41
<0.09
<0.11
2.46 ± 0.22
<0.38
<0.38
<0.38
<0.38
<0.24
77.50 ± 0.95
<0.15
0.32 ± 0.08
0.26 ± 0.08
<0.15
CFHT 2001
[Ne V] 3426
[O II] 3727
H I 3735
H I 3751
H I 3772
H I 3798
H I 3836
[Ne III] 3869
H I 3889
H I 3970
He II 4027
C III 4069
H I 4102
He II 4201
C II 4267
H I 4340
[O III] 4363
N III 4379
He I 4471
He II 4543
[Ar V] 4626
O IV 4632
C III 4650
C IV 4659
C IV 4659
He II 4686
[Ar IV] 4711
[Ne IV] 4715
[Ne IV] 4725
[Ar IV] 4740
Hβ 4861 100.00 ± 0.29 100.00 ± 0.26 100.00 ± 0.78 100.00 ± 1.59
O V 4930
<1.2
N V 4945
<1.2
[O III] 4959
<1.2
<0.99 ± 0.29
[O III] 50071.82 ± 0.28
[Fe VI] 5146
<10.0
He II 54115.8 ± 0.22
He I 5876
<0.8
[Fe VII] 6087
<0.8
Hα 6563 322.03 ± 0.15 306.38 ± 0.22
[N II] 6584
<1.6
[S II] 6716
<1.6
[S II] 6731
<1.6
[Ar V] 7006
<1.6
He I 7065
<1.6
[Ar III] 7136
<1.6
[O II] 7320
<1.6
[O II] 7330
<1.6
[O II] 7333
<1.6
?? 7408 0.82 ± 0.24
[Cl IV] 7530
<1.6
He II 7600
<1.6
O V 7611
<1.6
O IV 7713
<1.6
C IV 7726
<1.6
[Ar III] 7751
<1.6
[Cl IV] 8046
<1.6
CIII 8196
<1.6
He II 82371.76 ± 0.16
H P16 8502
<1.6
H P15 8545
<1.6
H P14 8598
<1.6
H P13 8665
<1.6
H P12 8750 1.29 ± 0.26
H P10 9014 1.7 :
[S III] 9069
<1.6
<1.55
4.94 ± 1.18
4.36 ± 1.30
<3.0
6.73 ± 0.533.71 ± 0.82
<1.2
<1.51.67 ± 0.48
2.37 ± 0.47
5.85 ± 0.58
9.06 ± 0.43
13.29 ± 0.44
7.81 ± 0.87
13.24 ± 0.79
<1.2
2.1 ::
<1.5
<1.5
25.11 ± 0.4225.49 ± 0.6219.79 ± 1.06
<1.2
<1.2
<1
<1
46.09 ± 0.3848.06 ± 0.4442.49 ± 0.80
<1.2
<1.2
<1.2
<1
<1
<1
2.97 ± 0.332.91 ± 0.35
<1.2
<1.2
<1.2
<1.2
<1.2
<1
<1
<1
<1
<1
75.13 ± 0.3179.02 ± 0.2879.06 ± 1.11
<1.2
<1.2
<1.2
<1.2
<1
<1
<1
<1
<0.8
<0.8
<0.28
<0.28
0.59 ± 0.13
2.53 ± 0.16
<0.28
0.69 ± 0.69
1.81 ± 0.79 2.16 ± 0.26
<0.8
6.49 ± 0.285.12 ± 0.31
<0.19
<0.8
<0.8
248.54 ± 2.63
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G. Stasi´ nska et al.: The chemical composition of TS01
2.2.2. Kitt Peak data
TS01 was observed at the 4m telescope of Kitt Peak National
Observatoryon1January2003.Thegratingusedwas KPC-10A,
andtheslit 2??×300??, withanorientationofPA = 44.7◦. Twoex-
posures of 600s were obtained. The data were reduced by L.J.,
employing the same procedure as for the SDSS spectrum, ex-
plained below.
2.2.3. SDSS and Kitt Peak data
The spectrum of TS01 appears in the data of the Sloan Digital
Sky Survey SDSS (http://www.sdss.org) under the name
0953-52411-1602. We present its analysis performed by L.J.
We separated the nebular emission from the stellar spectrum
and evaluated the reddening with as few free parameters as pos-
sible. We assumed the stellar spectrum to be that of a single
white dwarf (WD), hence neglectingthe possible contributionof
the companion; we considered three model WD spectra at tem-
peratures of 90, 120 and 150kK (the same as used for the CFHT
2003 spectrum).As for the nebular continuum,we computedthe
free-free and free-bound emissivities of H+and He++with the
chianti code (Landi et al. 2006), assuming an abundance ra-
tio He++/H+= 0.075 and an electronic temperature of 30kK
(Richer et al. 2002). We also retrieved the Hγ nebular emissivity
at this temperature from Storey & Hummer (1995).
First, we computed a model of the total (stellar+nebular)
spectrum around the Hγ line. In each of the WD spectral mod-
els, the Hγ line has a Voigt profile with a given equivalent width
(EW) Wwd, Gaussian width σwdand Lorentzian width awd. As
for the nebularemission, we computedthe EWWnebof the emis-
sion line with respect to the nebular continuum. Furthermore,
we assumed the real width of the nebular line to be much
smaller than the instrumental one, so its observed Gaussian and
Lorentzian widths, respectively σinstand ainst, are representative
of the instrumental PSF. Consequently, the observed widths of
the stellar line are σ∗= (σ2
normalized the local stellar+nebular continuum with the fit of a
slope on either side of the line. Finally, we let the central wave-
lengths of the stellar and nebular line, respectively λ∗and λneb,
be independentfrom each other. Using the data and assumptions
gathered, we fitted a consistent model on the observed spec-
trum around the line. Calling Vneb(λ) the profile of the nebular
line, V∗(λ) that of the star (both being normalized to an EW of
unity), and C∗the stellar contribution to the flux, the model can
be written as
wd+ σ2
inst)1/2and a∗= awd+ ainst. We
Fλ(λ) = C∗(1 − WwdV∗(λ − λ∗; σinst,ainst,σwd,awd))
+(1 −C∗) (1 + WnebVneb(λ − λneb; σinst,ainst)).
For all three WD modelatmospheres,we obtained verygoodfits
with no visible systematic residuals.
Given that the intrinsic Hα/Hβ nebular line ratio in TS01
is not merely the case B recombination value, and given the
additional problems with Hα (see Sect. 5.3.1), we cannot use
this ratio to evaluate the extinction. Hence, we used the contin-
uum to measure the latter. We first corrected the data for the
(1)
2The analysis by L.J. was done soon after our discovery of the spec-
trum in SDSS data release 2. In data release 6 (Adelman-McCarthy
et al. 2008), SDSS spectra were recalibrated, resulting in an increase
of about 30% of the fluxes of TS01. Line ratios remained unchanged.
Therefore, the analysis of L.J. remains valid. On the other hand, when-
ever we needed to consider the TS01 continuum in this paper, we used
the recalibrated spectrum.
Fig.5. Fully processed SDSS spectrum (dereddened, free of continuum
and stellar lines).
Fig.6. Fully processed Kitt Peak spectrum (dereddened, free of contin-
uum and stellar lines).
small foreground extinction (E(B − V) = 0.029) estimated by
Schlegel et al. (1998). Then, we removed most of the nebular
or stellar lines from the observations making use of a median
filter. Finally, adopting the SMC extinction law (Prévot et al.
1984; Bouchet et al. 1985) and comparing the filtered spec-
trum with the theoretical continuum, we evaluated the redden-
ing and corrected the data for it. The choice of the extinction
law was motivated by the low metallicity of TS01. The redden-
ing amounts obtained (additionally to the foreground one) are
E(B − V) = 0.033, 0.044 and 0.050 for the 90, 120 and 150kK
WD models, respectively.
The last processing of the data was the removal of the stellar
and nebular continua, to avoid the contamination of the nebu-
lar lines by the underlying stellar features. We used the fit of
the Hγ line to shift the theoretical stellar spectrum and nebular
continuum according to their evaluated radial velocities, con-
volved them by the average instrumental PSF and subtracted
them. Finally, we identified visually the detectable lines and
measured their fluxes. The fully processed SDSS spectrum with
the 90kK WD model spectrum removed is presented in Fig. 5,
while Fig. 6 shows the result of the processing of the Kitt Peak
spectrum. The choice of the WD model had a moderate impact
on the evaluation of the line fluxes, of order of 2% for most of
them and2σ in the worst case. The intensities of the SDSS spec-
trum and the Kitt Peak spectrum are listed in Table 2.
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A&A 511, A44 (2010)
1E-142E-14
1.5’’
NV 1240
sky
sky
NIV 1488
HeII 1640
CIV 1550
stellar continuum
Fig.7. 2D calibrated HST FUV spectrum of TS01.
2.3. Ultraviolet spectroscopy
2.3.1. HST STIS data
The HST STIS spectroscopic data correspond to the same pro-
posal (ID 9466) as the imaging data.
A 52??× 0.5??slit was used. It was oriented along the bright
jet-like emission of the nebula (PA 103◦), see Fig. 2.
Far UV observations The MAMA detector combined with
a G140L grating provided 2D spectra o8do03020, o8do03030,
o8do03040, o8do03050 and o8do03060, with 4675 s exposure
time each. The spectra cover a wavelength range from 1170
to 1700Å, with a resolving power of 1190 at the central wave-
length 1425Å. The 2D spectra show a bright blue stellar con-
tinuum and a few faint and extended emission lines from the
nebula. Calibrated 2D spectra were combined (after shifting be-
cause the spectroscopic observations were dithered) to produce
a 389.6 min spectrum. The resulting 2D spectrum is shown in
Fig. 7. The stellar spectrumshows good signal-to-noiseand stel-
lar and interstellar absorption are present. The analysis of the
stellar spectrum is presented in a companion paper (Tovmassian
et al., in prep.). Regarding nebular lines, the following ones are
detected: Nvλ1240,Niv]λ1488,Civ]λ1550andHeii]λ1640.
Selective absorption of resonance lines by the intervening inter-
stellar medium is treated in Sect. 3.4.
From the combined 2D spectrum the nebular emission was
extractedonbothsides ofthecentralstar,withanextractionwin-
dow of 60 pixels, equivalent to 1.464??. Figure 8 is a combina-
tion of both nebular spectra. The line fluxes in each lobe and the
combined values with respect Heii]λ1640 are listed in Table 3.
Near UV observations The MAMA detector combined with
a G230LL grating provided 2D spectra o8do02010, o8do02020,
o8do02030, o8do02040 and o8do02050, covering a wavelength
range from about 1600 to 3150Å. The calibrated 2D spectra
were combined (after aligning) to produce a spectrum with a
total exposure time of 237.5 min. As for the FUV, the NUV stel-
lar spectrumhas goodsignal-to-noise,andstellar and interstellar
absorption can be seen. However, no nebular lines are detected.
In particular, He ii λ1640, Niii]λ1750, and Ciii]λ1909 are not
seen.Table3givesupperlimitsforthelatterlineintensities,with
respect to He ii λ1640, as seen in the FUV spectrum.
2.3.2. FUSE data
The
Spectroscopic Explorer (FUSE) and their reductions (done
by G.T.) are described in Tovmassian et al. (2004). No emission
observationsof TS01with the
Far Ultraviolet
Fig.8. HST FUV nebular spectrum of TS01, not corrected for red-
dening, showing the Heiiλ1640Å, Nvλ1240Å, Niv]λ1488Å, and
Civλ1550Å lines. The fluxes are in 10−14ergcm−2s−1Å−1.
Table 3. Observed HST UV line fluxes, relative to He II 1640.
Ion lambda
Heiiλ1640Å
Nvλ1240Å
Oivλ1402Å
Niv]λ1488Å
Civλ1550Å
Niii]λ1750Å
Ciii]λ1909Å
Bright lobe
1.56a
0.31
<0.06
0.17
1.10
<0.1
<0.1
Faint lobe
1.15a
0.51
–
noisy
1.33
Combined
(2.7 +/− 0.3)a,b
0.47 +/− 0.04
< 0.06
0.12 +/− 0.05
1.28 +/− 0.10
<0.1
<0.1
Notes.(a)The flux of Heiiλ1640 is in units of 10−14erg cm−2s−1;(b)all
the flux in Heiiλ1640, including both lobes.
lines were detected in the observed wavelength region between
900 and 1200Å, except the H Lyβ line. For the photoionization
modellingof the nebula,it is importantto determineupperlimits
to the intensities of nebular lines expected in this wavelength
range. We proceeded in the following way. From a previous
model of TS01 we took the computed nebular continuum. We
superimposed on it the lines Ciiiλ977.020, Niiiλ989.799 and
Heiiλ992.4 with FWHM of 0.1Å (which corresponds to the
measured expansionvelocity of 30 kms−1, see Sect. 3.2).To this
we added the central star model mentioned in Sect. 2.2.1. The
resulting spectrum was processed through the interstellar hy-
drogen absorption simulator (http://violet.pha.jhu.edu/
~gak/fwebsim.html)to be compared with the observations. It
turned out that the corresponding lines start to be detectable in
the resulting spectrum when the line flux reaches approximately
7.5 × 10−14ergcm−2s−1. Indeed, the wavelength region that we
are exploiting here is very complicated. Apart from the different
interstellar absorptions and terrestrial airglow, the lines in this
region also lie at the edges of the detectors where they overlap,
and errors are much higher compared to other regions to the red.
2.4. Mid-infrared spectroscopy
TS01wasobservedusingtheinfraredspectrograph(IRS,Houck
et al. 2004) on board the Spitzer Space Telescope (Werner et al.
2004) on 22 April 2006 (program #20358). The observations
used the short-high (SH: 9.9−19.6 μm; R ∼ 600) and long-
high (LH:18.7−37.2μm; R ∼ 600) modules. The aperture of the
SH module is 4.7??× 11.3??and of the LH one is 11.1??× 22.3??,
so the entire nebular flux was measured. The details of the
Page 6 of 19

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G. Stasi´ nska et al.: The chemical composition of TS01
Fig.9. LH infrared spectrum of TS01 between 22 and 28μm, showing
the [Oiv]λ25.89μm and [Nev]λ24.32μm lines.
performed observations are shown in Table 1. For LH we used
four exposure cycles of 240 s each for on-source and off-source
observations, while for SH only on-source observations were
performed with six exposure cycles of 480 s each. The starting
points for our interactive data reduction were the co-added 2D
flat-fielded BCD (basic calibrated DATA) images (one for each
node position; pipeline version 15.3 for SH and 17.2 for LH).
The rogue pixels were removed using the IRSCLEAN tool3,
withtheaggressiveparameterequalto0.Thenthedatawerepro-
cessed (full extraction, trimming, defringing and averaging over
cycles) into a single spectrum per node position using SMART4
(Higdon et al. 2004). A similar procedure has been applied for
LH off-source observations, and the obtained spectra have been
subtracted from the on-source data for the corresponding node
position, to cancel out the sky background. The resulting spec-
trum between 22 and 28μm is shown in Fig. 9. For the high-
resolutionSH modulenobackgroundsubtractionwasdonesince
no sky measurements were taken and the SH slit is too small for
on-slit background subtraction. Finally, the spectra obtained for
both modules were averaged over two node positions, and the
detected nebular lines were measured within SMART.
The resulting intensities are listed in Table 4 together with
the estimated uncertainties. These uncertainties do not include
possible calibration errors. It is generally considered that the ab-
solute flux calibration has an accuracy of 20−30%. This will be
taken into account in the modelling (see Sect. 4.2).
Table 4 also lists the blue-shifts of the lines. One can see that
theyareconsistentwith theopticalmeasurementsofTovmassian
et al. (2001).
2.5. XMM data
TS01 has been also observed in the X-rays with XMM. The
data acquisition and analysis is presented in Tovmassian et al.
(in prep.).
3This tool is available from the Spitzer Science Center website:
http://ssc.spitzer.caltech.edu.
4SMART was developed by the IRS Team at Cornell University and
is available through the Spitzer Science Center at Caltech.
Table 4. Observed mid IR line fluxes, in units of 10−21W cm−2.
Ion lambda FluxUncertainty Rad. velocity
km s−1
–180
–140
10−21W cm−210−21W cm−2
0.38
1.50
0.82
[Oiv]λ25.89μm
[Nev]λ24.32μm
[Nev]λ14.32μm
0.03
0.20
0.10
3. What else do we know about TS01
and its exciting star?
3.1. Extinction
TS01 suffers only little extinction. Using the observed Hγ/Hβ
and Hδ/Hβ ratios, Richer et al. (2002) had found E(B − V) ∼
0.3 mag. However, this estimate was made without considering
the underlying stellar absorption in the Balmer lines. Due ac-
count for this effect significantlyreduces the estimated E(B−V),
as noted by Tovmassian et al. (2004). The extinction can also
be estimated by considering the spectral energy distribution of
the stellar core as observed in the far UV by FUSE. Assuming
a temperature of 120kK for the central star, Tovmassian et al.
(2004) obtained a good fit to these observations for E(B − V) =
0.045 mag, when using a non-canonical value for RVof 2.3 and
the interstellar reddening tables from Fitzpatrick (1999). Such a
low value of RVcompared to the standard one of 3.1 was con-
sidered compatible with the location of TS01 well outside the
galactic disk, since the intervening dust is likely composed of
smaller grains than in the spiral arms. However, we now know
that the temperature of the star which dominates the UV con-
tinuum is much cooler (see Sect. 3.3 and Tovmassian et al., in
prep.), implying that a steep reddening law is not needed af-
ter all. In the remainder of the paper as well as in Tovmassian
et al. (in prep.) we use the Fitzpatrick (1999) reddening law
parametrized with RV = 3.1, and take E(B − V) = 0.03 mag,
which satisfactorily accounts for the observed H and He line ra-
tios as well as the observed continuum. Note that the absence of
an absorption dip at 2200Å imposes an upper limit of 0.06 for
E(B − V).
3.2. Expansion velocity
The expansion velocity of TS01 has been measured by Richer
et al. (2003). This parameter is useful to estimate the expansion
cooling in the nebula. It also allows one to have an idea of the
nebular dynamical age. We adopt vexp= 30 km s−1.
3.3. The stellar core
Our understanding of the stellar core of TS01 has evolved
considerably since the first paper where it was suggested
that SBS1150+599A is a high excitation planetary nebula
(Tovmassian et al. 2001) with a central star that has an effec-
tive temperature above 100000K. Spectroscopic variations in
the course of one single night, reported in Tovmassian et al.
(2004), indicated the presence of a double system with a com-
pact star. Photometric observations then unambiguously deter-
mined a period of 3.92h (Napiwotzki et al. 2005). Analysis of
the light curve indicated that the visible star is likely an elon-
gated ellipsoid irradiated by a source of higher energy. It also
supported the previous conclusion that the companion must be a
(pre-?)white dwarf or a neutronstar. Finally, X-rayobservations
(Tovmassian et al. 2007, 2008) obtainedwith the XMM-Newton
Page 7 of 19

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A&A 511, A44 (2010)
satellite directly revealed the light from the companion, which
turns out to be a hot compact star! Thus, as will be shown later,
the “cool” star is the one visible in the optical and the UV and
it provides most of the ionizing photons. But it is the “hot” star
which gives rise to the high excitation lines observed in the neb-
ula. The best fit to the total spectral energy distribution of the
binary core indicates that the cool component has a temperature
Tc∼ 60kK, while the hot component should have Th∼ 170kK.
However,thedeterminationofthetemperatureofthehotcompo-
nent is not very accurate. Note that, in the scenario developped
by Tovmassian et al. (in prep.), the hot component is an old
white dwarf, which has a 170kK temperature not because it is
still early on its cooling path, but because it was heated by nu-
clear burningof the accreted material on its surface. For the cool
component Tovmassian et al. (in prep.) obtains the following:
Tc= 58000± 3000K, loggc∼ 5.1. The lower limit on the tem-
perature is the intrinsic temperature of the star, the upper limit
corresponds to the zone heated by irradiation. It is important to
note that the cool componentis not spherical and has not onlyan
inhomogeneous temperature distribution on its surface but also
an uneven gravitational acceleration. Its total luminosity is esti-
mated by Tovmassian et al. (in prep.) to be Lc = 1700L?with
about 30% uncertainty. Below we will consider for the sake of
simplicity that the cool star is sufficiently well represented by a
stellar modelatmospherewith Tc= 55kK and loggc∼ 5.1,with
a total luminosity of 1700L?.
The abundance analysis performed by T.R. on the cool star
gives 12 + logHe/H = 10.95 and 12 + logC/H = 7.20, with an
uncertainty of about 0.3dex, and upper limits 12 + logN/H <
6.92 and 12 + logO/H < 6.80.
3.4. Interstellar absorption of nebular UV lines
In the course of his stellar atmosphere analysis, T.R. noted that
the observed Civλ1550Å and Nvλ1240Å absorption lines
were stronger than predicted by his best models. He suggested
that these lines are probably affected by interstellar absorption.
In that case, the intensities of the Civλ1550Å and Nvλ1240Å
nebular lines are also affected by absorption. Since these lines
are crucial for the determination of the nebular abundances in
TS01, we here explain how we corrected for this effect.
We use the following notations (all quantities are a function
of wavelength): FS: flux extracted at the position of the star; FN:
flux extractedat the adjacent positionin the nebula; F∗: real stel-
lar flux; Fneb: real nebular flux; Fsky: sky emission and nebular
continuum.Theoptical depthdue to interstellarabsorptionis de-
noted τ. We have
FS= (F∗+ Fneb+ Fsky)exp(−τ)
and
FN= (Fneb+ Fsky)exp(−τ),
so that, in the spectrum analyzed by R.T., we have:
FS− FN= F∗exp(−τ).
Concerning the Civλ1550 line, reading out from Fig. 10 we
find F∗exp(−τ) = 8.7 × 10−15(black line in the figure), leading
to exp(−τ) = 0.75. Therefore, if the measured nebular flux is
Fnebexp(−τ) = 1.28 × 2.7 × 10−14= 3.45 × 10−14(last column
of Table 3), the nebular flux after correction for absorption is
Fneb= 4.65 × 10−14with an uncertainty of about 20%.
Concerning the Nvλ1240Å line, we find from Fig. 11 F∗=
2.1 × 10−14(red line in the figure), F∗exp(−τ) = 1.9 × 10−14
Page 8 of 19
-13.7
-13.6
-13.5
1170 1180
log fλ / erg/cm2/sec/A
o
wavelength / A
o
Teff = 50
55
60
kK
C IVC III N III N IV
-14.0
-13.9
-13.8
1540 15501560
C IV N V C IVN III
Fig.10. Stellar model atmosphere fitting of carbon lines in the stellar
core of TS01. Observation is in black.
-14.1
-14.0
-13.9
17101720 1730
log fλ / erg/cm2/sec/A
o
wavelength / A
o
Teff = 50
55
60
kK
N IV
-13.8
-13.7
-13.6
1230 12401250
C IV N V
O III
N V C III
Fig.11. Stellar model atmosphere fitting of nitrogen lines in the stellar
core of TS01. Observation is in black.
(black line in the figure), leading to exp(−τ) = 0.9. Therefore, if
the measured nebular flux is Fnebexp(−τ) = 0.47×2.7×10−14=
1.27 × 10−14(last column of Table 3), the nebular flux after cor-
rection for absorption is Fneb= 1.41 × 10−14. The uncertainty is
larger here, since the line is weaker. We adopt 30%.
4. Photoionization modelling
4.1. Global strategy
The chemical composition of TS01 can only be determined
throughphotoionizationmodelling,sincewehavenodirectelec-
tron temperature diagnostic. With the observational data now at
hand, we are able to confine the range of possible abundances
much better than in previous studies. In this paper, we try to
make the best use of all the observational constraints. The first
aspect concerns the morphology. The HST image (see Fig. 1)
has an elliptical shape, with two distinct narrow lobes. As the
Hα surface brightness distribution shows a hole in the centre, it
indicates that those lobes are real (possibly due to jets) and not a
thin disc seen edge-on. Because the presence of these lobes can
affect the ionization structure of the nebula, we choose to carry
out the photoionization modelling with a code that allows us to
deal with such geometries: Cloudy_3D (Morisset 2006), based
on Ferland’s 1D code CLOUDY (Ferland et al. 1998). We use
version c07.02.01 of CLOUDY and version 594 of Cloudy_3D.

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G. Stasi´ nska et al.: The chemical composition of TS01
Fig.12. Density structure of the nebula. Left: the chosen density struc-
ture of the model along the polar axis and along an axis perpendicular
to it; middle: the resulting Hα surface brightness distribution along the
same axes (continuous lines) compared with the observed distribution
(dotted lines); right: the theoretical Hα image.
We assume that the nebula is axisymmetric with its large
axis in the plane of the sky and that the lobes have a circular
cross-section. By trial and error,we choose a density structureto
reproduce the observed Hα surface brightness distribution. The
chosen density law along the polar axis and along an axis per-
pendiculartoitis shownintheleftpanelofFig.12.Theresulting
Hα surface brightness distribution along the same axes is shown
in the middle panel of this figure (continuous lines) and is com-
pared with the observed distribution5(dotted lines). The right
panel shows the theoretical Hα image, which can be compared
with the observed image shown in Fig. 1, especially as regards
the width of the polar lobes. Note that the density contrast be-
tween the lobes and the main body of the nebula is very modest:
only a factor of about 2. The density distribution is parametrized
by n0, the value of the density at the centre. For each run, we
choose n0in a way that, within a circle of a radius of 2.5??, our
model returns an Hβ flux of 2.5×10−14erg cm−2s−1, which cor-
responds to the observed extinction-correctedvalue from Richer
et al. (2002)6.The valueof n0is thusdependentonthe distanced
for which the computations are made.
The distance d in turn results from a fitting of the theoretical
optical/UV continuum to the observed one (taking into account
nebular continuum, aperture effects and reddening).
For the stars, we use model atmospheres computed by T.R.
with the Tubingen NLTE model atmosphere package (TMAP).
For the cool component, we use models tailored for our object.
For the hot component, in absence of sufficient observational
constraints, we chose among the complete flux tables for H-Ni
models with halo composition (May 2001) downloaded from
http://astro.uni-tuebingen.de/~rauch. Those models
are described in Rauch (2003).
4.2. The ultraviolet, optical and infrared fluxes on the same
scale
After a model has been run, the extinction-corrected line inten-
sities are computed for each of the observing slits and are com-
pared to the observations. This is the best way to deal with aper-
ture corrections, in particular when combining UV and optical,
or IR and optical data. Indeed, such a procedureaccounts for the
ionization structure of the object under study.
Absolute calibration of spectroscopic observations is noto-
riously difficult. We intercalibrate the UV/optical data by forc-
ing the measured value of the He ii λ1640/He ii λ4686 ratio
to the one predicted by our photoionization models in the cor-
responding slits. The value of f(STIS), representing the factor
by which the measured UV fluxes have to be multiplied for the
5For comparison with the model, we have symmetrised the observed
nebular surface brightness.
6Note that the models are not ionization bounded.
Heii λ1640/Heii λ4686ratio to bein agreementwith the model,
lies between 0.90 and 0.92 in our models. The value of f(STIS)
is higher for models with higher electron temperature. To allow
an easier comparison between models and observations, we fix
the value of f(STIS) to 0.91.
For Spitzer-IRS observations we multiply the observed
fluxes by a factor f(IRS) which adjusts the observed values of
[Nev]λ24.3μm/[Ne v] λ3426 (after reddening correction) to
the one predicted by the photoionization model in the corre-
spondingslit. The values of f(IRS) range between 0.87 and 0.95
for the models we considered. It might be judged unreasonable
to scale infrared fluxes using the [Nev]λ24.3μm/[Ne v] λ3426
ratio. However, in the electron temperature domain relevant for
TS01, this ratio does not vary very strongly (from Te= 20kK
to 40kK, it decreases by only a factor of two). In any case, this
is the only option we have to link the Spitzer line fluxes with the
optical ones, since our Spitzer data contain no H or He lines. Of
course we bear this difficulty in mind in the discussion. To re-
move the model dependance of the IR fluxes correction, we fix
the value of f(IRS) to 0.91.
The fact that both f(STIS) and f(IRS) are found very close
to unity is remarkable and means that the flux calibration of the
STIS and IRS LH spectra of TS01 is excellent.
4.3. Judging a model
To judge a model it is convenient to divide the line ratios to be
fitted into different categories:
– ratios of hydrogen lines or of helium lines: they probe the
reddening law, the stellar underlying absorption, and the re-
combination line theory;
– ratios of two different ions of the same element, such as
[Oiv]λ25.9μm/[O iii] λ5007, N v λ1240/[N iv] λ1486,
[Ne v] λ3426/[Ne iii] λ3869, and [Ne v] λ3426/[Ne iv]
λ4720. They basically test whether the ionization structure
is well reproduced by the model. In this category, we add
the He ii λ4686/Hβ line ratio, which is more dependent on
the ionization level of the nebula than on the abundance of
helium;
– ratios of lines used to determine the chemical composition:
[Oiv]λ25.9μm/Hβ, C iv λ1549/Hβ, N v λ1240/Hβ, and
[Ne v] λ3426/Hβ. We also consider [O iii] λ5007/Hβ (al-
thoughit is redundantwith [Oiv]λ25.9μm/Hβ oncethe ion-
ization structure is reproduced).
Note that in the case of TS01, the only ratio for direct plasma
diagnostics (i.e. electron temperature and density) that is avail-
able is [Nev]λ24.3μm/[Ne v] λ3426. Unfortunately, the two
lines come from measurements in different apertures and with
different observing techniques, and as mentioned above, there
is a priori some uncertainty in the relative calibration of the two
wavelengthdomains.However,thefactthatwefind f(IRS)close
to unity argues that the electron temperature of the [Ne v] emit-
ting zone in our models is not far from the true one.
For all the observables considered (usually line ratios), we
compute the value of
κ(O) = (logOmod− logOobs)/τ(O),
where Omodis the value returned by the model, Oobsis the ob-
served value, and τ(O) the accepted tolerance in dex for this
observable. For each observable, the value of τ(O) is chosen
a prioriconsideringthe observationalerrorbar,includingtheun-
certainty due to reddening,and the expectedability of ourmodel
(2)
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A&A 511, A44 (2010)
Fig.13. Graphic chart to compare the reference model with the mea-
sured line ratios of TS01. The numbers on the top of the panels indicate
the values of κ(O) for the line ratios listed at the bottom. The line ratios
presented in this figure were computed for the following pairs of slits
(using the same nomenclature as in Table 5): (0/0), (4/4), (3/0), (3/4),
(1/1), (0/0), and (0/0) for the ratios testing the degree of ionization, and
(0/0), (3/4), (1/0), (1/0), and (0/0) for the ratios used to determine the
abundances.
to reproduce a given observable. The value of τ(O) is defined as
follows:
τ(O) = log(1 + ΔO/O),
(3)
where ΔO is the absolute value of the maximum “acceptable”
error on the observable. We then judge our models by looking
at their outputs presented in graphical form (see an example in
Fig. 13). A model is fully satisfying only if each of the values of
κ(O) is found between −1 and +1, and if the computed line in-
tensities satisfy the upper limits for undetected lines. Of course,
a preliminaryconditionfor a model to be consideredis that it re-
turns the correct value of the Hβ flux, as explained in Sect. 4.1.
4.4. The reference model
Here we present our reference model, R. This is the model for
which all the values of κ(O) are as close as possible to zero, tak-
ing the following characteristics for the cool star: Tc = 55kK,
loggc= 5.1, Lc= 1700L?, and applying the extinction correc-
tions with RV = 3.1 and E(B − V) = 0.03 mag as explained in
Sect. 3.1. The model contains graphite grains (as expected for
a carbon-rich planetary nebula, which is the case of TS01 as
seen below). The grains have a standard size distribution and a
total dust-to-gas mass ratio of one tenth of the standard value.
A larger abundance of grains would bring the predicted contin-
uum around 24μm into conflict with the observation. Below we
explore deviations from the reference model which still account
for the observational data.
The reference model has Th = 170kK, loggh = 6.7, and a
total luminosity Lh= 2564L?7; n0= 181cm−3and the follow-
ing abundances,in unitsof 12+log(X/H):He = 10.95,C = 7.84,
N = 7.15, O = 6.82, Ne = 6.83, S = 5.65, Ar = 4.70.
Themodelsandobservationstowhichtheyarecomparedare
presentedin Table 5. Column1 of this table lists the line identifi-
cations, Col. 2 characterizes the observation using the following
nomenclature: 0 for CFHT 2003, 1 for STIS, 2 for FUSE, 3 for
Spitzer, 4 for SDSS, 5 for CFHT 2001. For lines which belong
to a wavelength range that was not observed, the number 6 is
attributed. Column 3 lists the observed reddening-correctedline
intensities (or their upper limits), in units of Hβ = 100 in the
corresponding apertures. Column 4 lists the acceptable relative
error ΔO/O used to computeκ(O). In the case of HST, FUSE and
Spitzer data,we estimatethe valueofHβ in therelevantaperture,
based on our models (since, as explained in Sect. 4, they deliver
a smoothed version of the observed surface brightness distribu-
tion).Thetoprows ofCol.5 ofthe tablelist the characteristicsof
the referencemodel.The predictedline intensities in the relevant
aperture are given in the following rows in units of Hβ = 100 in
the same aperture. For easier analysis, the next rows list a few
important line ratios, where the intensity of each line is mea-
sured through the aperture corresponding to the observation. In
order to shorten the table, we do not list the lines for which the
predictions from all our models give values smaller than 0.001
of Hβ (we note that this is the case for all the recombination
lines of elements C, N, O). We list only the strongest H and He
lines (we checked that the weaker H i and He ii lines always
give |κ(O)| < 1−1.5 in our models, implying that the correction
for stellar absorption and reddening is satisfactory).
The graphical representation of the line ratios predicted by
modelR andusedto estimate the chemicalcompositionofTS01
is shown in Fig. 13. This is the kind of chart that was used in
practice when judging the models that were run. A “best model”
is one for which all the diamonds fall as close as possible to
the ordinate 0. In any case, an acceptable model should have all
line ratios represented by a diamond between the two horizontal
lines, which represent a one-sigma deviation from the observed
value. In addition, acceptable models should not return line in-
tensities above the upper limits allowed by the observations.
Figure 14 shows the monochromatic images of the refer-
ence model in various emission lines. It reveals a few interest-
ing features of the model: some lines, such as [O iii] λ5007
and [Ne iii] λ3869 arise mainly in the lobes. Other lines, such
as C iv λ1549, N v λ1240, and [Ne v] λ3426 line come from
the entire nebula (in agreement with what Fig. 3 suggests for
[Nev] λ3426),while Ovi λ1032 and [Nevi] λ7.6μm (the latter
not in the observed wavelength range) come from the innermost
regions. We can also see that the two Ciii lines, Ciii] λ1909and
C iii λ977, although produced by the same ion, come from dif-
ferent regions: C iii λ977 has an important component coming
from the central main body (see Fig. 14), where the very high
electron temperature allows for its excitation even if C++is not
very abundant there.
Figure15 comparesthe
tion from TS01 with the computed one. The top panel
shows the reddening-corrected flux computed for model R,
in ergcm−2s−1Å−1, in the wavelength range of 900Å−40μm.
The observations are superimposed in various colours, as
indicated in the caption. One can see that the model reproduces
observedenergydistribu-
7The mass of the hot stellar component in the models depends on the
value assumed for the gravity, which is not well constrained. For the
photoionization modelling, what really matters is Thand Lh.
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G. Stasi´ nska et al.: The chemical composition of TS01
Fig.14. Monochromatic images of the reference model in various lines
(the values of the wavelengths are in Å if they are larger than 900, and
in μm otherwise). The x and y values are the coordinates in pixel units
of the models.
the observed spectral energy distribution quite well (except for
the IRS SH observations which could not be corrected for sky
emission, as explained in Sect. 2.4). The bottom panel shows
the energy distribution in the soft X-ray range: the blue curve is
the hot star, the green triangles are the XMM observations. The
flux from the star has been corrected for the nebular absorption
(computed by CLOUDY) and for the interstellar absorption,
taking a hydrogen column density of 1.6 × 1020cm−2. For each
computed model we checked that the ionizing flux does not
violate the observed stellar emission up to 200eV. At higher
energies, the observed emission may have another origin than
the stars we consider, but it does not affect our model fitting,
since we have no relevant observationalconstraints (the ion with
the highest ionization potential observed is Ne4+, which has an
ionization potential of 97.1eV).
Figure 16 compares the energy distributions of the two stars
considered in the modelling: the “cool” star is in red, the “hot”
one is in blue. The sum of the two is in black. As mentioned in
Sect. 3.3, the cool star dominates in the optical range (a few eV)
and until about 20 eV, but it is the hot star which provides the
photons with energies above 40−50eV. Consequently it is the
cool star which providesmost of the H ionizingphotons,but it is
the hot star which provides the photons responsible for the pres-
ence of the He ii, N v, [O iv], [Ne iv], [Ne v] and [Ar v] lines.
This is averyuncommonsituation,perhapsa uniquecase among
planetary nebulae: TS01 has two ionizing stars! This explains
why our previous attempts to model the object were facing the
difficulty that the nebula needed plenty of photons of energies
Fig.15. Comparison of the reddened spectrum of the reference model
with observations. Top: from the UV to the IR. The reddened model is
in black. The colour code for the observations is as follows. Magenta:
FUSE; blue: HST; green: SDSS ; red: Spitzer (the SH observations
could not be sky-corrected). Bottom: X-ray domain. The model (with
extinction applied) is in blue. The XMM observations are represented
by triangles.
Fig.16. Spectral energy distribution (normalized to arbitrary units) of
the radiation from the ionizing stars in the reference model. Red: the
“cold” star; blue: the “hot” star; black: the sum of the two.
above 54.4 eV, while the Balmer absorption lines in the stellar
continuum indicated a moderate temperature.
From Figs. 13, 15, and Table 5 one can see that our refer-
ence model fits all the observational constraints very well. The
only exception is that of Ciii] λ1909, whose intensity is slightly
abovethe upperlimit we gaveto the STISobservation.However,
we consider this result to be still acceptable, since upper limit
fluxes for unobserved lines are difficult to estimate accurately.
The abundances of C, N, O, and Ne in the reference model are,
respectively, 1/3.5, 1/4.2, 1/70, and 1/11, and, for S and Ar <
1/30 of the solar values given by Asplund et al. (2005).
5. The chemical composition of TS01
5.1. Range of abundances for the reference model
We now investigate the error bars on abundances that are due
only to the uncertainties in the observed line intensities. Since
C, N, O, and Ne contribute very little to the energy bud-
get, it is straightforward to estimate from model R the mini-
mum and maximum abundances corresponding to the minimum
and maximum values of [Oiv]λ25.9μm/Hβ, C iv λ1549/Hβ,
Page 11 of 19

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A&A 511, A44 (2010)
Table 5. Photoionization models versus observations. Line intensities are in units of Hβ = 100 in the corresponding aperture.
(1)(2)a
(3)(4) (5)
R
170
2564
6.7
181
(6)
Mi
170
1618
6.7
181
(7)
Ma
170
4064
6.7
181
(8)(9)
HeMi
170
1618
6.7
181
HeMa
170
4064
Th 103K
Lh
log gh
n0 cm−3
[L?]
6.7
181
Sunb
10.93
8.39
7.78
8.66
7.85
7.14
6.18
Hec
Cc
Nc
Oc
Nec
Sc
Arc
10.95
7.84
7.15
6.82
6.83
5.65
4.7
10.95
7.64
10.95
8.05
7.32
7.13
6.9
5.83
4.92
10.98
7.64
10.91
8.05
7.32
7.13
6.9
5.83
4.92
77
6.63
6.76
5.5
4.5
6.63
6.76
5.5
4.5
d [kpc]
Fβd
M(H) [M?]
f(STIS)
f(IRS)
22.5
2.55
22.3
2.52
0.14
0.91
0.88
22.9
2.56
0.15
0.91
0.94
22.3
2.54
0.14
0.92
0.87
22.9
2.54
0.15
0.91
0.95
0
0.91
0.91
ΔO/O
0.15
0.15
0.1
0.1
0.07
0.07
0.05
0.05
0.04
0.04
H8
H8
H7
H7
H6
H6
Hγ
Hγ
Hα
Hα
0
4
0
4
0
4
0
4
5
4
5
4
1
0
4
4
9.51
9.29
12.99
13.6
25.66
25.71
47.84
46.73
251.41
311.03
<0.19
<1.44
609.18
77.88
74.93
5.43
9.01
9.15
14.18
14.39
23.64
23.87
47.40
47.64
282.06
279.89
0.16
0.04
594.53
76.53
76.16
6.19
9.02
9.19
14.22
14.46
23.70
23.96
47.47
47.71
282.48
280.14
0.25
0.06
572.27
73.85
74.00
6.02
9.02
9.14
14.18
14.36
23.64
23.85
47.34
47.59
281.30
279.23
0.11
0.03
623.3
79.67
78.91
6.38
8.96
9.12
14.12
14.36
23.54
23.80
47.41
47.65
283.21
280.85
0.27
0.06
607.55
78.35
78.52
6.40
9.09
9.21
14.28
14.47
23.81
24.03
47.41
47.66
280.50
278.43
0.10
0.02
575.72
73.60
72.93
5.88
HeI5876
HeI7065
HeII1640
HeII4686
HeII4686
HeII5412
0.1
0.04
0.04
0.1
CIII]1909
CIII977
CIV1549
CIV4659
1
2
1
0
<61.44
<430.3
1054.84
<0.38
73.41
104.58
1050.89
0.14
89.01
112.54
890.87
0.08
59.96
96.78
1246.07
0.26
93.05
117.74
916.75
0.08
56.73
92.22
1200.76
0.26
0.2
[NIII]57.2
NIII]1750
NIII991
NIV]1486
NV1240
6
1
2
1
1
1.23
3.00
7.54
65.89
330.38
1.54
4.29
9.52
67.40
253.7
0.96
1.99
5.91
61.95
425.38
1.56
4.49
9.81
69.93
259.25
0.93
1.87
5.69
58.82
410.86
<60.84
<423.77
73.05
339.3
0.5
0.3
[OIII]4363
[OIII]5007
OIV]1402
[OIV]25.9
OVI1032
0
0
1
3
2
<0.42
2.52
<36.97
13.72
<407.42
0.11
2.57
7.68
13.57
50.48
0.14
3.24
6.88
11.20
30.45
0.12
2.58
10.60
20.35
105.06
0.15
3.35
7.15
11.24
30.65
0.11
2.49
10.05
20.13
102.45
0.3
0.3
[NeIII]3869
[NeIII]15.5
[NeIII]36.0
[NeIV]2424
[NeV]3426
[NeV]14.3
[NeV]24.3
[NeVI]7.6
0
3
3
6
0
3
3
6
0.79
<7.22
<28.88
0.50.82
0.21
0.02
25.20
58.32
43.25
54.44
27.47
1.45
0.37
0.03
29.46
48.63
35.05
44.16
15.98
0.44
0.11
0.01
20.51
66.92
50.69
63.76
45.07
1.50
0.38
0.03
30.10
49.19
34.84
43.90
15.84
0.43
0.11
0.01
19.93
65.84
50.91
64.03
45.48
57.67
29.62
54.17
0.2
0.3
0.2
[ArIV]4711
[ArIV]4740
[ArV]7005
0
0
4
<0.15
<0.15
<1.44
0.15
0.12
0.27
0.15
0.12
0.23
0.14
0.11
0.31
0.16
0.13
0.23
0.13
0.10
0.30
Page 12 of 19

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G. Stasi´ nska et al.: The chemical composition of TS01
Table 5. continued.
(1)(2)a
(3) (4) (5)
R
0.29
0.45
(6)
Mi
0.22
0.33
(7)
Ma
0.37
0.57
(8)(9)
HeMi
0.22
0.33
HeMa
0.36
0.57
[ArV]8.0
[ArV]13.1
6
3
<7.23
[SIII]9069
[SIII]18.7
[SIII]33.5
[SIV]10.5
4
3
3
3
<1.41
<7.22
<36.1
<7.23
0.24
0.21
0.34
6.65
0.31
0.25
0.41
6.29
0.17
0.17
0.28
6.88
0.32
0.25
0.42
6.35
0.16
0.16
0.27
6.74
NV1240/ NIV]1486
[OIV]25.9/[OIII]5007
[OIV]25.9/[OIII]5007
[NeV]3426/[NeIII]3869
[NeV]3426/[NeIV]4720
[NeV]24.3/ [NeV]3426
[NeV]24.3/ [NeV]14.3
1/1
3/0
3/4
0/0
0/0
3/0
3/3
4.645 0.7
0.139 0.4
0.194 0.4
73.021 0.7
99.039 0.3
0.024 0.3
1.828 0.3
5.01
0.13
0.22
71.33
127.42
0.024
1.26
3.76
0.087
0.136
33.47
89.66
0.023
1.26
6.86
0.20
0.35
3.71
0.085
0.134
32.84
87.90
0.020
1.26
6.98
0.21
0.36
151.80
181.34
0.024
1.26
154.26
185.72
0.025
1.26
Notes.(a)Labels for the observing slits: 0: CFHT 2003; 1: STIS; 2: FUSE; 3: Spitzer; 4: SDSS; 5: CFHT 2001;(b)Asplund, Grevesse & Sauval
(2005);(c)abundances in units of 12 + log X/H;(d)in units of 10−14erg cm−2s−1.
N v λ1240/Hβ, and [Ne v] λ3426/Hβ without changing the
ionization structure of the nebula. However, one has also to
consider the error bars on line ratios that constrain the ioniza-
tion structure: [Oiv]λ25.9μm/[O iii] λ5007, N v λ1240/[N iv]
λ1486, [Ne v] λ3426/[Ne iii] λ3869, and [Ne v] λ3426/[Neiv]
λ4720. The minimum values of the C, N, O, and Ne abun-
dances in TS01 are obtained by a model with the lowest ion-
ization compatible with the observations and the lower limits
of [Oiv]λ25.9μm/Hβ, C iv λ1549/Hβ, N v λ1240/Hβ, and
[Ne v] λ3426/Hβ. Such a model, Mi, is reported in Col. 6 of
Table 5. It is derived from the reference model R by lowering
the values of Lh, and decreasing the values of the C, N, O, and
Ne abundances.With similar considerations,one can construct a
model Ma, which will give the maximum C, N, O, and Ne abun-
dances. This model, with a higher Lhand same This listed in
Col. 7 of Table 5.
The resulting limits on the abundances of C, N, O, and Ne in
the gaseous phase8of TS01 are thus:
7.64 < 12 + logC/H < 8.05
7.00 < 12 + logN/H < 7.32
6.63 < 12 + logO/H < 7.13
6.76 < 12 + logNe/H < 6.90.
The limits on the C/O, N/O and Ne/O ratios are obtained by
consideringtailored models which reproducethe extreme values
of the observed intensities of the C, N and Ne lines. They are:
0.83 < logC/O < 1.21
0.12 < logN/O < 0.54
−0.16 < logNe/O < 0.17.
To derive the lower limit on He/H, one must consider the
model with the lowest ionization compatible with the ob-
served [Oiv]λ25.9μm/[O iii] λ5007, N v λ1240/[Niv] λ1486,
[Ne v] λ3426/[Neiii] λ3869, and [Ne v] λ3426/[Neiv] λ4720,
8The contribution of grains to the abundance of carbon is discussed in
Sect. 5.2.2.
and the lowest He ii λ4686/Hβ. The upper limit on He/H is ob-
tained with similar arguments. The correspondingmodels HeMi
and HeMa, respectively, are listed in Cols. 8 and 9 of Table 5.
The resulting limits for He/H obtained in this way are 0.095
and 0.081. In other words, the precision on the He/H abun-
dance is poor, despite the fact that we have been fitting the
He ii λ4686/Hβ ratio within 4% (the formal uncertainty in this
ratiois2%forthe2003CFHT data,butacomparisonwithSDSS
dataledustoadoptahighervalueforthetolerance).Withamore
accurateupperlimit onHeiλ5876/Hβ we couldreducethe error
bar on the helium abundance (and actually, model HeMi slightly
violates the present upper limit on He i λ5876/Hβ). But the un-
certainty in He/H will remain larger than the uncertainty in the
He ii λ4686/Hβ ratio mainly because the electron temperature
gradient in this nebula is steep and the ratio of emissivities of
He ii λ4686 and Hβ slightly varies with temperature.
5.2. Additional sources of abundance uncertainties
In this section, we discuss how reasonable variations of the pa-
rameters that were so far fixed in the modelling procedure af-
fect the derived abundances.We also discuss some more general
problems that may have an influence on the estimated chemical
composition of TS01. To save space, the models that were con-
structed to discuss these additional uncertainties are not listed
in the paper. We will only mention their impact on the derived
abundances.Note that all of those additionalmodels have the re-
quired angular size and total Hβ flux, and their abundances have
been chosen to fit the observed emission line ratios.
5.2.1. The effect of changing the description of the stars
So far, we have kept the parameters of the cool star fixed. Even
if they are rather well determined, as explained in Sect. 3.3 and
Tovmassian et al. (in prep.), it is important to see the effect that
a change in those parameters implies on the derived chemical
composition of the nebula. It turns out that an increase of 5kK
in Tcinduces a decrease in the C, N, O and Ne abundances by
0.05−0.08dex. A change in gcby 0.1dex leaves the abundances
of the fitted nebular model unchanged.
Page 13 of 19

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A&A 511, A44 (2010)
We then explored the effects of changing the parameters of
the hot star. By increasing the temperature of the hot star by
10kK, (implying a slight decrease of its luminosity to fit the
observations) one decreases the C, N, O, Ne abundances of the
fitted nebular model by about 0.2dex.
We also explored the effect of changing the model atmo-
sphere of the hot star. One extreme case is to consider a model
atmosphere composed only of H and He, instead of the Galactic
halo chemical composition. Consequently, there are no absorp-
tion edges in the atmosphere above 54eV. The model which fits
the observations presents an intense [Ne vi] λ7.6μm emission
(unfortunatelyoutside the wavelength range of our IRS observa-
tions of TS01). Its Ne abundance is higher by 0.15dex than that
of the reference model, while the abundances of the remaining
elements are almost unchanged.
5.2.2. Dust issues
Concerning the extinction and reddening issues, a change of
E(B − V) and RV within limits compatible with the observed
Balmer decrement and the observed stellar energy distribution
does not alter the abundances derived for TS01 significantly.
Our reference model has a dust-to-gas mass ratio
of 10−1times the canonical value, with the canonical grain size
distribution as stated in Sect. 4.4. The chemical composition of
the grains – pure graphite – is dictated by the fact that the object
is undoubtedly carbon-rich, as seen in Sects. 3.3 and 5.1. The
total abundance of grains in the reference model is chosen in a
way that the predicted infraredflux which arises from the heated
grains does not exceed the observed IRS LH flux and that it pro-
ducesnosignificantdipat 2200Å,sincethisis notobserved.The
total amount of carbon locked in grains in the referencemodel is
0.4 times that of the abundance of carbon in the gas phase. This
means that the total abundance of carbon in the nebula (gas plus
grains) is larger by about 0.15dex than given in Sect. 5.1.
5.2.3. The role of morphology
While constructing our reference model (and all the models de-
scribed before), we have chosen a geometry that reproduces the
observed Hα surface brightness, including the lobes. It is inter-
esting to experiment with a simpler model without any lobes,
in which the averaged surface brightness is the same as in the
reference model. The abundances in such a model differ in-
significantly from those of the reference model. We have to con-
fess that we were somewhat surprised by this result, since as
shown in Fig. 14 the emission in such lines as [O iii] λ5007
or [Ne iii] λ3869 traces the lobes very distinctly. On the other
hand, one has to remember that the density contrast between the
lobes and the ambient medium is only a factor of two, as seen
in Fig. 12.
5.3. Caveats
5.3.1. The problem of Hα
One of the intriguing problems in the observations of TS01 is
the behaviour of the Hα line. As seen in Table 2, the observed
Hα/Hβ ratio varies among data taken during different runs and
at different telescopes. Since the ratios of all the remaining hy-
drogen lines look normal within the error bars, we are inclined
to think that this Hα problem has no influence on the derived
chemical composition. Nevertheless, we feel it important to try
to understand the reason for the observed values of Hα/Hβ.
In the present study, we have done the computationswith the
full treatment of hydrogenas offered by CLOUDY (this, and not
case B, is actually the default option in CLOUDY). Under the
physicalconditionsinthisnebula,oneindeeddoesnotexpectthe
Balmer lines to be emitted undercase B, not even with the added
effect of collisional excitation. The ionization parameter of the
emitting regions is high and the nebula is optically thin, which
renders it a good candidate for case C as described by Baker &
Menzel (1938) and reconsidered by Ferland (1999). In such a
case, absorption of Lyman photons from the star contributes to
the emission of the Balmer lines, and the Balmer decrement de-
pends on the number of respective Lyman line photons in the
star. However, we are far from reproducing the Hα/Hβ ratios
observed in the various slits. Of course, the computed Balmer
decrement strongly depends on the fluxes at the wavelengths
of the H Lyman lines in the model atmosphere used. But the
differences in the Hα/Hβ ratios in the different observing runs
make it doubtful that simple stellar fluorescence can explain the
observations.
The reference model predicts a ratio of about 2.81. The dif-
ferences in the observedHα/Hβ ratios cannot have a nebular ori-
gin since the associated time scales are far too long.
Water vapour absorption near Hα is far too weak to explain
the variations9. Now that the nature of the binary central star
is better known, we can also discard the possibility that much
of the Hα emission comes from an accretion disc. Active mass
transfer in the system has ceased and, even if there is a stellar
wind or weak remnant of an accretion disc around the hot com-
ponent, it cannot have a big influence on emission lines, since
we detect fairly symmetric underlying absorption lines from the
cool componentat all orbital phases. These symmetric lines also
imply that extra emission from the irradiated face of the cool
component does not contribute any significant Hα emission.
The remaining option is atmospheric refraction (Filippenko
1982), since the slit was not oriented at the parallactic angle for
many (though not all) of the spectra with Hα/Hβ ratios differ-
ing significantly from 2.81 (rather, usually east-west). What is
odd a priori, if atmospheric refraction is responsible, is that the
lines from Hδ to Hβ are observed with constant intensity ratios.
Simulations in which we convolve the quantum efficiency of
the slit camera used at SPM10with the object’s very blue spec-
trum indicate that the effective wavelength is between 4000 Å
and 4500 Å. Thus, the effective wavelength, which is what is
used to centre the object in the slit, is between the blue lines,
so atmospheric refraction has very little effect upon them. As a
result, Hα should be the only optical line that may be signifi-
cantly affected by atmospheric refraction. Also, comparedto the
usual assumptions, the wavelength baseline over which atmo-
spheric refraction operates is unusually large in this case, of the
order of 2000Å or more. Tests using the SPM4 dataset (Richer
et al. 2002), in which this issue can be studied in greatest detail,
clearly implicate the effect of atmospheric refraction since the
spectral shapeofthecentralstar’s continuumvariesas a function
of the differencebetween the slit positionangle and the parallac-
tic angle. Therefore we are inclined to attribute the variations
observed in the Hα/Hβ ratio to atmospheric refraction.
9http://www.astrossp.unam.mx/sitio/
abs_telurica_english.htm
10http://www.astrossp.unam.mx//Instruments/bchivens/
camrend/manual-english.pdf
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G. Stasi´ nska et al.: The chemical composition of TS01
5.3.2. Atomic data
As noted by Péquignot & Tsamis (2005), the atomic data
on which photoionization models are built are not of per-
fect accuracy. All the models we have computed rely on
CLOUDY c07.02.01. It is not excluded that future advances in
atomic physics, especially in the calculation of recombination
coefficients for highly ionized species, might affect the com-
puted ionization structure. However, the fact that we now have
observational data (or stringent upper limits) on several ions of
each of the elements C, N, O and Ne in TS01 makes us confi-
dent in the robustness of the chemical composition that we have
derived. The relatively large error bars we obtain on the abun-
dances (principally due to the lack of a direct measure of the
electron temperature in the nebula) imply that the uncertainties
in atomic data, including the collision strengths of the lines used
for the diagnostics,shouldbe negligiblein the total error budget.
5.3.3. Dynamical effects
Schönberner et al. (2005) have drawn attention to the possible
importance of dynamical effects in the thermal balance of neb-
ulae. They make the point that the role of dynamical expan-
sion in the cooling budget increases as the metallicity decreases.
We have therefore included the effect of expansion cooling in
CLOUDY byintroducinga windcoolingcontributionin the rou-
tine ‘‘cool\_eval.cpp’’:
dynamics.dDensityDT = (float)(2.*fudge(0));
CoolHeavy.expans =
dense.pden*phycon.te*BOLTZMANN*dynamics.
dDensityDT;
with the user defined parameter“fudge” related to the expansion
velocity and the outer radius of the nebula by fudge”= vexp/Rout.
All the models presented above have been computed with
an expansion velocity of 30 km s−1, corresponding to the ob-
served value (see Sect. 3.2). We have tried other values for vexp
in the equation above, but noted no significant changes in the
output between 0 and 200 km s−1, the extreme values we tried.
This result is at variance with the finding by Schönberner et al.
(2005) that expansion cooling significantly reduces the temper-
ature with respect to a fully static model of the same density
structure.
In our models, the dominant cooling process is collisional
excitation of H Lyα, and, at the ionization level predicted by
the model, it is clear that expansion cooling must be negligible,
unless the velocity of the jet is of the order of 1000 km s−1.
Could it be that the lower temperaturefoundby Schönberner
et al. (2005) in fully dynamical models with respect to hydro-
static ones, which they attribute to expansioncooling,is actually
the result of some other process? The only idea that comes to
mind is a departure from ionization equilibrium. For an average
temperature of 30kK and an average density of about 200 cm−3,
the recombination time for hydrogen is about 103yr. From the
apparent size, expansion velocity and distance to TS01, one can
estimate an expansion time of ∼7 × 103yr. Therefore, the neb-
ula should not be far from ionization equilibrium. On the other
hand, the dynamical model shown in Schönberner et al. (2005)
was for a 0.595M?star with an effective temperature of 100kK,
corresponding to an evolution time of ∼5 × 103yr. The average
density of the nebula in their simulation is then about 100 cm−3.
In such a situation, the nebula is farther from ionization equilib-
rium. Since their star is in a phase where the number of ionizing
Table 6. Nebular abundances of TS 01, in various units.
12+log X/HUncertainty
±0.04
±0.30
±0.25
±0.33
±0.30
X/H Mass fraction
2.63 × 10−1
6.11 × 10−4
1.46 × 10−4
7.79 × 10−5
9.96 × 10−5
<7.45 × 10−6
<8.38 × 10−7
He
C
N
O
Ne
S
Ar
10.95
7.84
7.15
6.82
6.83
<5.5
<4.5
8.91 × 10−2
6.92 × 10−5
1.41 × 10−5
6.61 × 10−6
6.76 × 10−6
<3.16 × 10−7
<3.16 × 10−8
photons increases with time, the ionization level of the dynami-
cal model should be lower than that of the corresponding static
model. Hence, Lyman alpha cooling should be more important
andtheelectrontemperaturelowerthaninthehydrostaticmodel,
which is indeed what their dynamical model yields. In TS01
the dynamical effects on the ionization and temperature of the
nebula should be much smaller than in the case computed by
Schönberner et al., if noticeable at all. In their model, the tem-
perature drop due to dynamical effects is about 10kK. Given the
argumentation above, we consider that any dynamical effect on
the electron temperature in TS01 would be of 2−3kK at most
with respect to the temperature we compute in our model. As an
experiment, we computed a model where we use the CLOUDY
parameter“cextra”withavalueof10−20.3ergcm3s−1tosimulate
an extra cooling factor that reduces the average electron temper-
ature by about3kK with respect to the referencemodel. We then
adjusted the abundances to reproduce the observed line ratios.
We found that the abundances in this model are not very differ-
ent from those of model R. In particular, the abundance of O is
not changed. The reason is that the model must reproduce the
[Ne v] λ3426/[Neiii] λ3869 ratio, which is nearly temperature-
independent, and that ratios used to constrain the oxygen abun-
dance ([O iv] λ25.9 μm/Hβ and [O iii] λ5007/Hβ) are not very
sensitive to the temperature above 30kK.
5.4. Wrapping up
In summary,consideringall the possible sources of uncertainties
and adding in quadraturethe various independenterrors, we find
that the elemental abundances in the gas phase of TS01 are as
listed in Table 6.
An additional amount of carbon, about 40% of the total el-
emental abundance, is locked up in dust grains. Allowance for
this component raises the carbon abundance in the nebula to
12 + logC/H = 8.00 ± 0.3.
The error bars on the derived abundances may seem large
when compared to the typical error bars in other PNe. However,
one must remember that the analysis of TS01 is much more dif-
ficult, due to the absence of direct temperature diagnostics and
to the weakness of the lines from metals.
The abundances derived for the nebula are consistent with
those derived by T. R. for the atmosphere of the cool star, see
Sect. 3.3,exceptforcarbonwhoseabundanceis largerby0.8dex
in the nebula. This agreement is remarkable, given the difficulty
of the analysis. Whether the discrepancy between the carbon
stellar and nebular abundances is real should be examined in
more detail.
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