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Appendix A: Details on the SMARTS data reduction

Table A.1:   Comparison stars for the SMARTS photometry.

The ANDICAM optical detector is a Fairchild 447  CCD with 30 m pixels. It was used in  binning mode, yielding a nominal plate scale of 0.369 arcsec pixel^-1 and approximately a field. Readnoise was 6.5 e^- rms; gain is 2.3 e^-/DN. It was read with a single amplifier. Filters were Johnson B and V and Kron-Cousins R and I. We generally took 3 exposures in B and V, and 2 images in and , and co-added these prior to analysis. The IR channel detector is a Rockwell 1024² HgCdTe ``Hawaii'' array with 18 m pixels binned to yield a  pixel^-1plate scale. The IR channel field is . The used filters are standard CIT/CTIO JHK. We flattened the images using dome flats obtained every 3 nights. We observed using a 3-point dither pattern and a throw, and used the median of the unshifted images as the local sky.

A.1 Differential photometry

We initially determined relative magnitudes using differential photometry. In the optical, we selected the 11 brightest stars in the field of view (Table A.1; V372 Ori was overexposed). Instrumental magnitudes were determined for the target and all the comparison stars, by summing the counts in a 5 pixel ( ) aperture. Background was the median level in an annulus from 10-20 pixel ( ) radius. Since V1118 Ori is in a region of complex nebulosity, we also tried modeling the background by taking radial cuts through the image and extrapolating the level at the position of the target. This affected the magnitudes at less than 5% at B, and even less at longer wavelengths.

In the near-infrared, we used 6 comparison stars in the field (again excluding V372 Ori). We used a 5 pixel ( ) aperture. Background was the median level in an annulus from 10-20 pixel ( ) radius.

We constructed a single comparison, the ``superstar'' from a weighted sum of all the comparison stars. The differential magnitude was the difference between the instrumental magnitudes of the target and the ``superstar''.

Most of the stars in this field are known or suspected variables. Fortunately, none were as variable as V1118 Ori. We assumed that, by summing all the comparisons, the ``superstar'' would be less variable than any single comparison. We constructed light curves using subsets of the comparison stars, and did not see any gross effects that could be attributed to variations of the comparisons. Nonetheless we did not incorporate those stars with measured standard deviations after absolute calibration (see below) >0.10 mag at V into the optical ``superstar''. We caution that there is no certainty that the comparison was stable at the few percent level.

A.2 Optical absolute calibration

Although the SMARTS operations are not optimized for absolute photometry, single observations are made of Landolt standard fields on each night judged photometric to establish thephotometric zero point. These can be used to determine apparent magnitudes with an accuracy of better than 10%.

On each night on which we had both observations of the target and of a photometric standard field, we computed the zero-point offset (apparent magnitude - instrumental magnitude). We assumed the standard atmospheric extinction correction for Cerro Tololo. Since the standard field was observed only once per night, we could not account for changes in transparency through the night. We did not solve for reddening terms, so the photometric solution consisted of only the zero-point.

We used the zero-point correction to determine the apparent magnitudes of all the comparison stars in the field. We could do this on 71 nights. In doing so, we confirmed that all the comparison stars were variable (Table A.1). We compared these magnitudes to the magnitudes published by Hillenbrand (1997). The median offsets were 0.08 and 0.23 mag at V and , respectively. Given that all the stars are variable members of the Orion Nebula population, this was acceptable agreement.

We applied this absolute calibration on all nights (including non-photometric nights) to determine the apparent magnitudes of the target. There were 14 nights where SMARTS photometry overlapped with the Villanova photometry. On these nights there were systematic offsets, in the sense (SMARTS-Villanova), of  , , and  mag in V, , and , respectively. The cause for the systematic offsets might be the lack of a color term in the photometric solution, exacerbated by the fact that most of our comparisons have colors >2 mag. To allow comparison between the different telescope data, we applied the above systematic offsets to the SMARTS magnitudes and gave the nightly average data in Table A.2.

A.3 Near-infrared absolute calibration

For the JHK absolute calibration, we compared the instrumental magnitudes directly to the 2MASS magnitudes of the 6 comparison stars. To check for target variability, we compared these magnitudes to those published in the DENIS catalog, and by Ali & Depoy (1995) and Hillenbrand et al. (1998). Two of the six comparisons (2MASS J05344159-0534249 and 2MASS J05344219-0533036 appear variable, with variances >0.2 mag at K(4 observations), and 2 others (JW 94 and JW 100) appeared variable in the optical (Table A.1).

Our ability to absolutely calibrate the photometry was ultimately limited by the stability of the zero-point. By using all the possible comparison stars, on average the zero-point will be more stable than determined from any single comparison star. We did not include the systematic uncertainty in the zero-point in our error budgets. Overall, the JHK fluxes compared well to those published by Lorenzetti et al. (2007). However, we added a systematic correction of +0.1 mag for J and +0.2 mag for K to match their photometry. V1118 Ori is intrinsically variable, as are the comparison stars. The only reason for applying the systematic offsets to our data was to simplify graphical comparison with the Lorenzetti et al. data. The offsets do not change significantly the results of this paper, they merely shift the near-infrared data points by 0.1 to 0.2 mag, e.g., in Figs. 10 and 11.

Table A.2:   SMARTS average nightly magnitudes.

Table A.3:   Villanova magnitudes.

Appendix B: Details on the Spitzer data reduction

B.1 IRAC

We used the MOsaicker and Point source EXtractor (MOPEX) software release of June 2007. We started from the Basic Calibrated Data (BCD) individual images to produce mosaics with a (native) pixel size of 1 22. For the data taken in sub-array mode (PIDs 3716 and 41019), we collapsed the 3-dimensional image BCDs into 2-dimensional BCDs by using, for each pixel, the median of the 64 planes. For the uncertainty BCD files, we used the standard deviation of the 64 planes of the BCD file instead of using the input uncertainties. This method allows the removal of most particle hits and the ingestion of the collapsed BCDs in MOPEX. We then used aperture photometry, centered on V1118 Ori, using an extraction radius of 10 pixels (12 2), and a concentric annulus of radii 12 and 15 pixels (we used radii of 10 and 12.5 pixels for the IRAC4 band at 8.0 m because of the strength of the background intensity at this wavelength for radii larger than 12.5 pixels. Indeed, the background is dominated by emission caused by the nearby Herbig Ae star V372 Ori; this effect is less prominent in the other IRAC bands, but the presence of the Herbig star limited us to outer radii of 15 pixels). No aperture correction is needed for an extraction circle radius of 10 pixels (according to the IRAC data handbook v3.0).

B.2 MIPS

We used the BCD files where V1118 Ori was on the detector as input files for the MOPEX pipeline. A native pixel size of was used to create the mosaic, and we used an extraction radius of 13 centered on V1118 Ori to derive the aperture photometry. We used an annulus of radii 15 and 20 for the background, and finally used an aperture correction factor of 1.17 (from Table 3.12 of the MIPS data handbook, v. 3.3.1). We also investigated the background with a different method, i.e., to calculate the background in a nearby region using a circle with radius 13 . Indeed, the MIPS24 background near V1118 Ori is highly inhomogeneous. Notice that, while the uncertainty in the flux density is on the order 0.5 mJy, we estimate the true uncertainty, mainly because of the inhomogeneity of the background and the difficulty of finding a representative background region, to be closer to 15-20 mJy.

B.3 IRS

We extracted the IRS spectra using the Spitzer IRS Custom Extraction (SPICE) v2.1.2 software and the post-BCD co-added 2-D images (pipeline version S15.3.0 for SL/SH and S17.2.0 for LH/LL for the outburst data; S18.5.0 for all post-outburst spectra). We investigated in detail the methodology for subtracting the background in the IRS spectra. Indeed, for the low-resolution modules, we have investigated different techniques for estimating the background contamination. Firstly, we used the standard technique that takes the image from nod A, subtracts the nod B image, and then the net (A-B) spectrum of V1118 Ori is extracted with SPICE. The same is then done for the net (B-A) spectrum, and both net spectra are averaged. Unfortunately, this technique is unsatisfactory because of the inhomogeneity of the background emission along the cross-dispersion axis of the slit. A second technique consisted in using the images taken for the SL2 nods and use the source-free region in the SL1 image for the background of the SL1 nods. This method removes any contamination by V1118 Ori in the SL1 image, but the downside is that the region of the sky covered by the SL1 slit (during the SL2 observations) is away from V1118 Ori. Thus, the measured background may not represent the background near V1118 Ori. A third technique consisted in using a background region near V1118 Ori obtained during the observed nod. This technique allows to get a better estimate for the background near V1118 Ori, but has the disadvantage that V1118 Ori may contaminate the background, especially at longer wavelengths, since the standard extraction width increases with increasing wavelength, to take into account the increasing size of the point spread function of a point source. For the February 2005 observation, technique #2 proved better for SL1, while technique #3 was better for SL2. We used the PAH emission at 6.2 and 11.3 m (coming mostly from the diffuse background emission) to check that the background emission was adequately removed. For the March 2005 observation, we used technique #3 for both SL slits, as technique #2 gave an excess in the PAH at 11.3 m in the SL1 slit. For the post-outburst observations, we used technique #3 for both SL slits and also for the LL slits (in this case, the background subtraction is more difficult for  mic).

	Figure B.1: Spitzer IRS high-resolution post-outburst spectra (thick curve) degraded to the resolution of the SL and LL modules. The latter data are shown as well (thin curve). Note the good agreement in the continuum and for most lines, except for [Ne II].
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Figure B.2: Spitzer MIPS 24 m image taken in March 2004 with overlays of the IRS modules during the observations on 2005 February 18 ( top) and on 2008 November 14 after the outburst ( bottom). We only show the first position of the two nods, for clarity. Both nods for the background SH/LH observations are shown in the bottom figure. A scale is shown on both North-orientated images. The image scale is linear from 0 to 250 MJy sr-1.

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For the high-resolution modules, we used the standard approach to average the spectra from both nods using the full slit aperture. For the post-outburst spectra, we used the accompanying background spectra to obtain the background-subtracted spectra. This procedure worked nicely for the SH module (i.e., up to 19.35 m), as demonstrated by the good subtraction of the PAH emission at 11.3 m. Notice that the H₂ 0-0 S(1)  m line was apparently completely subtracted, while the H₂ 0-0 S(2)  m line is faintly detected, We also detect [Ne II]  m ( ²P_1/2-²P_3/2) and [S III]  m ( ³P₂-³P₁) in the background-subtracted spectrum (we will come back to this issue below). Furthermore, the silicate feature and continuum flux up to 14 m are consistent with the SL flux, suggesting that the background subtraction was not too far off. For the LH module, we subtracted only 95% of the background flux for the post-outburst data. Indeed, at the wavelengths covered by the LH, the background dominates the emission and is highly position-dependent. Although our background pointing was close to V1118 Ori, there is considerable inhomogeneity in the diffuse emission of the Orion nebula, as observed in the 8.0 and 24 m images. Figure B.2 shows the MIPS 24 m images together with the ``field-of-views'' of the IRS slits during the first outburst observation and in late 2008. We used the [Si II]  m ( ²P_3/2-²P_1/2) as the proxy for background subtraction for the LH module, and noticed that a scaling factor of 0.95 was better than 1.0. We detect, however, excess flux in the [S III]  m ( ³P₁-³P₀) line and faint excess in the H₂ 0-0 S(0)  m line (in contrast to the S(1) line in the SH spectrum). The SH and LH spectra agree relatively well at 19.4 m, suggesting again that the background subtraction was adequate.

We provide, however, a word of caution about the accuracy of the continuum flux level (especially above 30 m) and in the reality of the detected flux excess in some lines, especially in the LH wavelength region. Indeed, we have scaled the background LH spectrum using the [Si II] line. However, we remind the reader that V1118 Ori is located at the center of the Orion nebula, where strong ionization and excitation of the interstellar medium takes place. Strong spatial variations of the diffuse background emission occur, which may be different for the continuum and line (mostly [Si II], [S III], [Ne II], and H₂) emissions. Therefore, using a specific emission line to check that the background subtraction is adequate may result in an incorrect subtraction of the continuum emission of the nebula and even in the subtraction of another emission line! Nevertheless, the similarities of the continuum shape in the low-resolution and high-resolution modules (at least below 35 m) for the post-outburst data, the presence of line excess in [S III] in the LL spectrum and in both SH and LH spectra are suggestions that the background subtraction was overall accurate. The presence of [Ne II] in the SH module but not in the SL module demands, however, deeper analysis to confirm the detection of the line. We degraded the SH and LH spectra down to SL and LL resolution (Fig. B.1). The continuum shapes are well matched, but the degraded SH data covering the [Ne II] line shows a strong line which is not compatible with the SL data at this wavelength. This strongly suggests that the SH background [Ne II] line flux was too faint compared to the line flux near V1118 Ori and that the measured SH line flux is of background origin (from the Orion nebula or from diffuse emission close to V1118 Ori). In the case of H₂ at 12.28 m, the degraded SH spectrum is consistent with the SL spectrum: the contrast between the line and the continuum is too faint to confirm the presence of this line in the SL spectrum. The same comment applies to the other molecular hydrogen line at 28.2 m. Thus, we cannot confirm that this molecule is detected in V1118 Ori's spectrum. On the other hand the [S III] lines in the degraded SH and LH spectra have similar peak flux densities as in the LL spectrum, indicating that these lines originate from the immediate vicinity of V1118 Ori. In summary, we believe in the detection of the [S III] lines, while we have doubts for [Ne II] and the H₂ lines, but we nevertheless provide the line fluxes from the high-resolution module spectra in Table 3. In any case, in view of Spitzer spatial resolution, higher spatial resolution observations would be required to confirm the origin of the excess line emission.

We have further investigated whether the post-outburst background observations could be used for the outburst SH and LH data. In principle, the zodiacal light dominates the continuum background emission below about 40 m, while the interstellar medium emission dominates above. The zodiacal contribution varies as a function of time. We have used the Spitzer Planning Observations Tool (SPOT) to determine the contribution at the three different epochs of the IRS observations and found that they were of similar level (about 20-25 MJy sr^-1 from 15 to 24 m). For the SH spectrum, we preferred to use the post-outburst spectrum without applying a scaling factor, since the PAH feature and the H₂ line S(2) were well subtracted (the S(1) line is slightly oversubtracted). For the LH spectrum, we used a scaling factor equivalent to 0.85, i.e., about the value of the ratio of the zodiacal light contribution at 24 m at in February-March 2005 and November 2008. This ratio also cancels out relatively well the [Si II] line while the H₂ S(0) line still remains detected. With this procedure, the resulting spectrum shows an increase of the flux with longer wavelengths, perhaps coming from an envelope, but this increase is not detected in the post-outburst data, casting some doubt on the accuracy of the continuum shape in the SH and LH outburst spectra above 14 m. Since we have no outburst LL data to confirm this shape, we prefer to err on the safe side and consider that the high-resolution continuum shape for  m during the outburst is unreliable. The strengths of the emission lines may also be affected by unreliable background reduction; indeed the flux levels are generally larger than in the post-outburst spectrum (except for [Ne II], but in this case, the line is probably not originating in the direct vicinity of V1118 Ori). Line fluxes from the high-resolution module data in outburst are also given in Table 3.