EDP Sciences
Free Access
Volume 507, Number 3, December I 2009
Page(s) 1517 - 1530
Section Interstellar and circumstellar matter
DOI https://doi.org/10.1051/0004-6361/200912413
Published online 08 October 2009

Online Material

Appendix A: Presence of an X-ray component

The UV spectra exhibit a strong stellar O/ VI 1031, 1037 Åline (Fig. D.2) that cannot be reproduced by a $\sim$40 kK stellar atmosphere model. This line is classically explained by the Auger ionization due to X-ray emission produced by wind instabilities. A weak X-ray emission is confirmed by Chandra observations[*]. The Chandra data reduction has been performed using the Ciao V3.4 package CALDB (updated in May 2007) following standard procedures. Images in the 0.5-1.5 keV, 1.5-2.5 keV, and 2.5-8 keV bands were generated using ``dmcopy'' and ``csmooth'' Ciao procedures. Counts are only significantly detected in the softer energy band image (0.5-1.5 keV), indicating that the source is intrinsically soft.

The limited count detection (less than 40 photons) does not allow a robust spectral analysis. The soft energy flux (i.e. 0.5-2 keV band) is estimated by assuming an emission model of a black-body with a temperature of $\sim$0.03 keV and taking into account a neutral absorption of $N_{\rm H}=2\times 10^{21}~\rm {cm}^{-2}$. The unabsorbed measured flux of the source in the 0.5-2 keV band is $1 \times 10^{-14}~\rm {erg/s/cm}^2$.

To take this X-ray emission into account, we add a black-body emission at 300 kK to the stellar flux. We do not have any constraints on the shape of the X-ray SED emission, so we used the Planck function as the simplest SED.

This soft X-ray component does not affect the determination of the nebular parameters, such as the abundances. It does have a small influence on the high ionization potential (IP) emission lines ([Ne/ III] and [Ar/ IV] ), which are slightly underestimated by the model if the X-ray component is not taken into account. We adjust the luminosity of the soft X-ray component, so that these high IP emission lines are reproduced. The required bolometric luminosity is 3 $L_\odot$ (using 1.25 kpc for the distance), which corresponds to a flux in the 0.5-2 keV band of $4 \times 10^{-17}\rm {erg/s/cm}^2$, much lower than the observed X-ray flux. A hotter X-ray source would have a lower luminosity to conserve the soft X-ray flux and no influence on these lines. In contrast, a cooler X-ray source would require a higher luminosity, leading to intensities for these high IP lines that are higher than the observed values.

Appendix B: Line profiles

Gesicki et al. (1996) published observations of high spectral resolution of IC 418 for the /, [O/ III] 5007 Å, and [N/ II] 6584 Ålines. Sharpee et al. (2004) also presented high resolution observations for various ions. In both cases, the H I profiles are broad, with half width at half maximum (HWHM) close to 18 km s-1, the [O III] lines are narrower, with a HWHM close to 9 km s-1, while the low density species, such as [N II], show a double-peaked profile, with an inter-peak spacing of 20 km s-1 and each peak having a HWHM comparable to the one of the single [O III] line.

We compute the line profiles and compare them to the observations of Gesicki et al. (1996) using the same aperture size and position they used, namely a centered square aperture of 3 arcsec size. Despite the density law being determined by these authors in no way reproduces the observed surface brightness, we can fit an expansion velocity law similar to the one that they adopt. The resulting profiles reproduce the main properties of the observed profiles (see Fig. B.1). The velocity law we adopt here is $V \propto R^4$ with a maximum of 30 km s-1 at the outer edge of the nebula. It is beyond the scope of this paper to make a complete dynamical model of IC 418. We only check that the morphology we used is not trivially ruled out by the line profile observations. The values for the velocities of O++ and N+ that we found are compatible with the radiation-hydrodynamic models presented by Schönberner et al. (2005).

\end{figure} Figure B.1:

Emission-line profiles: models (lines) and observations (Gesicki et al. 1996, symbols), for / (black solid line and diamonds), [O/ III] 5007 Å(blue dot-dashed line and triangles), and [N/ II] 6584 Å(red dashed line and squares). Intensities are scaled so that the maximum of each profile reaches 1.0.

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Appendix C: Chlorine collision strengths and electron temperature diagnostic

The collision strengths for Cl++ used in Cloudy are taken from Wilson & Bell (2002). Note that Cloudy takes into account the correction by a factor of 10 for the collision strength for 1D2-1S0 as discussed in Keenan et al. (2003). We test these values against the previous values from Krueger & Czyzak (1970) and found that these latest are in closer agreement with the observations. On one side, the values for [Cl II] and [Cl III] are fitted when using values from Krueger & Czyzak, while it is impossible to fit both ions with the same Cl abundance when using collision strengths from Wilson & Bell. In this case, the [Cl II] lines are predicted to be too high, while the [Cl III] lines are fitted well.

On the other hand, the electron temperature diagnostic [Cl/ II] 6162/9124 Åis found to be predicted coherent with the [S II] and one of the two [O II] diagnostic when using the data of Krueger & Czyzak. For these low ionization lines, the model overpredicts the ratios with $\kappa([\mbox{S {\sc ii}}], [\mbox{Cl {\sc ii}}] ) \sim 1.5$. Using the values from Wilson & Bell leads to a prediction of the [Cl/ II] 6162/9124 Åratio that is lower than what is observed, with $\kappa([\mbox{Cl {\sc ii}}]) \sim -1.2$.

Appendix D: Online figures and tables

\end{figure} Figure D.1:

The fits to the nitrogen and FeIV lines. Note the nebular emission line C/ II 1761 Å.

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\end{figure} Figure D.2:

The fits to the oxygen lines. Argon and some FeV lines are also indicated.

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\end{figure} Figure D.3:

The fits to the carbon lines. N/ IV 1183, 85 Åand N/ V 1188 Ålines are also marked.

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\end{figure} Figure D.4:

Phosphorus and silicon lines. Note the strong absorption lines.

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\par\includegraphics[width=8.cm]{12413F17} \includegraphics[width=8.cm]{12413F18}
\end{figure} Figure D.5:

Quality factors $\kappa ({\rm O})$ versus the wavelength ( upper panel) and critical density ( bottom panel).

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\end{figure} Figure D.6:

Same as Fig. 6, but for a black-body model.

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\end{figure} Figure D.7:

Same as Fig. 7, but for black-body model.

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\par\includegraphics[width=9cm]{12413F21} \end{figure} Figure D.8:

Nebular abundances by mass from the literature (small crosses, same as in Table 2) and from the adopted model (large stars represent the stellar abundance, and large diamond the nebular abundances). Units are in log(abundance by number) relative to the solar values from Asplund et al. (2005).

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Table D.1:   Diagnostic line ratios, see text for the columns definitions.

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