A multi-wavelength study of the young star V1118 Orionis in outburst
M. Audard, G.S. Stringfellow, M. Güdel, S.L. Skinner, F.M. Walter, E.F. Guinan, R.T. Hamilton, K.R. Briggs, C. Baldovin-Saavedra
aa r X i v : . [ a s t r o - ph . S R ] D ec Astronomy&Astrophysicsmanuscript no. aa c (cid:13)
ESO 2018November 9, 2018
A multi-wavelength study of the young star V1118 Orionis inoutburst ⋆ M. Audard , , G. S. Stringfellow , M. G ¨udel , S. L. Skinner , F. M. Walter , E. F. Guinan , R. T. Hamilton , ,K. R. Briggs , C. Baldovin-Saavedra , ISDC Data Center for Astrophysics, University of Geneva, Ch. d’Ecogia 16, CH-1290 Versoix, Switzerland Observatoire de Gen`eve, University of Geneva, Ch. des Maillettes 51, 1290 Versoix, Switzerland Center for Astrophysics and Space Astronomy, University of Colorado, Boulder, CO 80309-0389, USA Institut f¨ur Astronomie, ETH Z¨urich, 8093 Z¨urich, Switzerland Department of Physics and Astronomy, Stony Brook University, Stony Brook, NY 11794-3800, USA Department of Astronomy and Astrophysics, Villanova University, Villanova 19085, PA, USA Department of Astronomy, New Mexico State University, 320 East Union Ave, Apt. 1434, Las Cruces, NM 88001, USAReceived 2009 July 31; accepted 2009 December 16
ABSTRACT
Context.
The accretion history of low-mass young stars is not smooth but shows spikes of accretion that can last from months andyears to decades and centuries.
Aims.
Observations of young stars in outbursts can help us understand the temporal evolution of accreting stars and the interplaybetween the accretion disk and the stellar magnetosphere.
Methods.
The young late-type star V1118 Orionis was in outburst from 2005 to 2006. We followed the outburst with optical andnear-infrared photometry; the X-ray emission was further probed with observations taken with
XMM-Newton and
Chandra duringand after the outburst. In addition, we obtained mid-infrared photometry and spectroscopy with
Spitzer at the peak of the outburst andin the post-outburst phase.
Results.
The spectral energy distribution of V1118 Ori varied significantly over the course of the outburst. The optical flux showedthe largest variations, most likely due to enhanced emission by a hot spot. The latter dominated the optical and near-infrared emissionat the peak of the outburst, while the disk emission dominated in the mid-infrared. The emission silicate feature in V1118 Ori isflat and does not vary in shape, but was slightly brighter at the peak of the outburst compared to the post-outburst spectrum. TheX-ray flux correlated with the optical and infrared fluxes, indicating that accretion affected the magnetically active corona and thestellar magnetosphere. The thermal structure of the corona was variable with some indication of a cooling of the coronal temperaturein the early phase of the outburst with a gradual return to normal values. Color–color diagrams in the optical and infrared showedvariations during the outburst, with no obvious signature of reddening due to circumstellar matter. Using Monte-Carlo realizations ofstar+disk+hotspot models to fit the spectral energy distributions in “quiescence” and at the peak of the outburst, we determined thatthe mass accretion rate varied from about . × − M ⊙ yr − to . × − M ⊙ yr − ; in addition the fractional area of the hotspotincreased significantly as well. Conclusions.
The multi-wavelength study of the V1118 Ori outburst helped us to understand the variations in spectral energy distri-butions and demonstrated the interplay between the disk and the stellar magnetosphere in a young, strongly accreting star.
Key words. accretion, accretion disks – Infrared: stars – stars: circumstellar matter – stars: coronae – stars: pre-main-sequence –stars: individual (V1118 Ori) – X-rays: stars
1. Introduction
The star formation process involves strong accretion of circum-stellar matter onto the protostar. The time evolution of the massaccretion rate is of deep interest to understand the timescaleof stellar growth and lifetime of proto-planetary disks. Whilethe mass accretion rate in young stars overall decreases withincreasing stellar age (Hartmann et al. 1998), it can also showsignificant and rapid changes over time (e.g., Hartmann et al.1993). A handful of accreting young stars display powerful erup-tive events with flux increases in the optical regime of a fewmagnitudes. Two classes have emerged: FUors, which displayoutbursts of 4 magnitudes and more, last several decades and,
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[email protected] ⋆ therefore, show a low recurrence rate. EXors (named after theprototype EX Lup), in contrast, show somewhat smaller out-bursts ( ∆ V = 2 − mag) on a much shorter timescale, froma few months to a few years, and may occur repeatedly (seereview by Hartmann & Kenyon 1996; also Herbig 2008). Suchoutbursts are believed to originate during a rapid increase of thedisk accretion rate over a short period of time, from values of − M ⊙ yr − to − M ⊙ yr − , although the underlying causeof such increase in mass accretion rate is unclear (thermal diskinstabilities, Lin & Papaloizou 1985; Bell & Lin 1994; closecompanions, Bonnell & Bastien 1992; cluster-induced encoun-ters, Pfalzner et al. 2008; giant planets in the disk, Clarke & Syer1996; Lodato & Clarke 2004; a combination of gravitational in-stability and the triggering of the magnetorotational instability,Armitage et al. 2001; accretion of clumps in a gravitationallyunstable disk, Vorobyov & Basu 2005, 2006). The limited num-ber of eruptive young stars and the long recurrence time (espe- M. Audard et al.: The 2005 Outburst of V1118 Ori cially for FUor-type objects) make it difficult to test models. Itis, therefore, important to study in as much detail as possible theevolution of outbursts and to understand the place of outburst-ing young stars in the evolutionary scheme from highly accret-ing, embedded protostars to “normal” accreting classical T Tauristars (CTTS).The optically revealed CTTS can make excellent compar-ison stars to outbursting young stars: they have been studiedextensively in the past: optical and infrared observations havegiven clues about the presence of accretion disks, on the massaccretion rates, and disk winds, while millimeter observationsprovided constraints on the dust disk mass and on CO outflows.In the X-rays, the origin of emission in young, accreting stars hasgained significant attention in the last few years. Kastner et al.(2002) obtained the first X-ray grating spectrum of TW Hya, amoderately accreting CTTS, and observed unexpected new spec-tral features: the O
VII and Ne IX He-like triplets showed verylow forbidden-to-intercombination line ratios and indicated highelectron densities of the order of − cm − . The effectof UV photoexcitation on the observed line ratio was deemednegligible. In addition, the X-ray spectrum was consistentwith a quasi-isothermal plasma of about 3 MK. Kastner et al.(2002) concluded that the X-ray emission in TW Hya was dueto accretion, a result further supported by Stelzer & Schmitt(2004) and Ness & Schmitt (2005). Additional high-resolutionX-ray spectra of CTTS also indicated that accretion may play asignificant role for X-rays in accreting stars (Schmitt et al. 2005;G¨unther et al. 2006; Robrade & Schmitt 2006; Argiroffi et al.2007; G¨unther et al. 2007; Sacco et al. 2008). Soft X-rayemission from shocks in jets may also be detected (G¨udel et al.2005, 2008; Kastner et al. 2005; Robrade & Schmitt 2007;Schneider & Schmitt 2008; G¨unther & Schmitt 2009).Telleschi et al. (2007) and G¨udel & Telleschi (2007) showedthat CTTS display a soft X-ray excess from plasma at lowtemperature. Such a plasma (mostly detected in the O VIIlines) is difficult to detect with CCD spectroscopy but caneasily be revealed with high-resolution X-ray spectra when theabsorbing column density is not too high. The soft X-ray excesscan co-exist with the hot plasma observed in the vast majorityof young stars, which is due to scaled-up solar-like magneticactivity (Smith et al. 2005; Audard et al. 2005a; Preibisch et al.2005; Telleschi et al. 2007). The origin of such an excess isunclear, but it could be due in part to accretion onto the stellarphotosphere. It is crucial to understand how the X-ray emissionin young stars, in particular those accreting matter actively, canbe influenced by the accretion process. FUor and EXor stars are,therefore, the ideal cases to understand the physical mechanismcommon to both CTTS and outbursting stars.In the recent years, several young low-mass stars wereobserved in outburst, in the optical and infrared, but alsoin the X-rays: V1647 Ori in 2003–2005 (e.g., Kastner et al.2004, Grosso et al. 2005, Kastner et al. 2006a; see alsoAspin et al. 2008 and references therein) and recently in 2008–2009 (Itagaki 2008; Aspin 2008; Venkat & Anandarao 2008;Aspin et al. 2009), V1118 Ori in 2005-2006 (Audard et al.2005b; Lorenzetti et al. 2006, 2007; Herbig 2008). Kastner et al.(2006b) also observed with Chandra a young stellar object ineruption in LDN 1415 but failed to detect it. EX Lup also showedan extreme outburst in early 2008 (Jones 2008a,b; Kospal et al.2008). ´Abrah´am et al. (2009) detected crystalline features in thesilicate feature of EX Lup in outburst which were not present inthe pre-outburst spectrum, suggesting crystallization by thermalannealing in the surface layer of the inner disk. X-ray obser-vations were obtained with
Chandra (PI: Weintraub) and
Swift (PI: Stringfellow) but are not yet published. We note also that theFUor objects FU Ori, V1735 Cyg, and Z CMa were also detectedin X-rays (Skinner et al. 2006, 2009a; Stelzer et al. 2009).The V1647 Ori campaign showed an increase of the X-rayflux up to a factor of 200 from its pre-outburst flux, in line withthe flux increase in the infrared. The X-ray flux then followedthe optical outburst flux and returned to its pre-outburst levelafter the outburst ended (Kastner et al. 2006a). The initial ob-servations of V1118 Ori in X-rays showed a different behavior,with little flux enhancement (Audard et al. 2005b), but also in-dicated that the rapid increase of accretion rate in outbursts canimpact the X-ray emission of young accreting stars. The presentpaper aims to present the remainder of the X-ray data of V1118Ori taken during and after the 2005–2006 outburst, together withcontemporaneous optical and infrared data.
2. The young erupting star V1118 Ori
The outburst of V1118 Ori, a young low-mass M1e star in theOrion Nebula ( d = 400 pc; Muench et al. 2008 for a discus-sion), was reported in early January 2005 by Williams et al.(2005). Hillenbrand (1997) and Stassun et al. (1999) providedetails about its physical properties: M ⋆ = 0 . M ⊙ ; R ⋆ =1 . R ⊙ ; P rot = 2 . ± . d; L bol ≥ . L ⊙ ; log T eff [ K ] =3 . ; log t [ yr ] = 6 . . Recently, Reipurth et al. (2007) re-solved V1118 Ori into a close binary separated by . ′′ witha position angle of ◦ and a magnitude difference of ∆ m =0 . mag in the H α -band image. The observation was obtainedon 2004 Jan 29, thus before the 2005–2006 outburst (see Fig. 6).The binarity of V1118 Ori leaves unclear which star actuallyerupted. However, we note that the small magnitude differenceindicates that the components are similar, suggesting similar ef-fective temperatures. We may also assume that they show similardisk evolution. While the binarity of V1118 Ori complicates theinterpretation of the combined photometry and spectroscopy inquiescence, this should have little impact on the interpretationof the outburst data, as only one star+disk component dominatesthe emission. In this paper, we have used the above stellar prop-erties that assumed a single star.V1118 Ori has shown frequent outbursts (e.g., 1983-84,1988-90, 1992-94, 1997-98; see Garcia & Parsamian 2000 andHerbig 2008 for details). In fact, after returning in quiescencein mid-2006, V1118 Ori had another outburst in late 2007(Garcia & Parsamian 2008). We focus here on our monitoringcampaign to study the 2005–2006 outburst of V1118 Ori in theX-rays, optical and infrared. Additional properties in the opti-cal and infrared were obtained independently during and afterthe outburst by Lorenzetti et al. (2006, 2007) and Herbig (2008).Using wind models, Lorenzetti et al. (2006) derived a mass lossrate of × − M ⊙ yr − from the H I recombination line, and (3 − × − M ⊙ yr − from the CO emission at 2.3 µ m for theneutral molecular gas; they also found no evidence of infraredcooling from a collimated jet or outflow. Lorenzetti et al. (2007)also found evidence of intrinsic polarization in the I band, and ir-regular fluctuations during the outburst. Herbig (2008) obtainedKeck/HIRES spectra of V1118 Ori in the decaying phase ofthe outburst and in the post-outburst phase. He noted the de-tection of Li I λ in emission during the outburst, in contrastwith the absorbed feature in CTTS spectra. The feature, how-ever, returned in absorption after the outburst. A similar behav-ior occurred in the K I lines at λλ α during the outburst, which disap-peared thereafter. In the near-infrared, Lorenzetti et al. (2007)also found that the emission lines detected during the outburst . Audard et al.: The 2005 Outburst of V1118 Ori 3 Table 1.
X-ray observation log.
Parameter Sep 2002 Jan 2005 Feb 2005 Mar 2005 Sep 2005Satellite . . . . . . . . . . . . . . . . . . . . .
Chandra Chandra XMM-Newton XMM-Newton XMM-Newton
ObsID . . . . . . . . . . . . . . . . . . . . . . 2548 6204 0212480301 0212480401 0212481101Duration (ks) . . . . . . . . . . . . . . . . 48 5 20 20 20Observation date . . . . . . . . . . . . . 2002 Sep 6–7 2005 Jan 26 2005 Feb 18-19 2005 Mar 21 2005 Sep 8UT . . . . . . . . . . . . . . . . . . . . . . . . . 12:57–02:54 03:07–05:04 22:36–04:50 16:20–22:08 12:43–18:51Average JD - 2,450,000 . . . . . . . 2,524.3 3,396.7 3,420.6 3,451.3 3,622.2Jan 2006 Feb 2006 Mar 2006 Apr 2006 Jul 2006 Dec 2007Satellite . . . . . . . . . . . . . . . . . . . . .
Chandra Chandra XMM-Newton Chandra Chandra Chandra
ObsID . . . . . . . . . . . . . . . . . . . . . . 6416 6417 0403200101 6418 6419 8936Duration (ks) . . . . . . . . . . . . . . . . 30 30 100 30 30 37.5Observation date . . . . . . . . . . . . . 2006 Jan 4–5 2006 Feb 23 2006 Mar 2–3 2006 Apr 23 2006 Jul 24 2007 Dec 14–15UT . . . . . . . . . . . . . . . . . . . . . . . . . 18:33–03:09 09:41–17:47 19:02–21:06 14:42–23:31 07:50–16:49 14:49–01:13Average JD - 2,450,000 . . . . . . . 3,740.5 3,790.2 3,797.8 3,849.3 3,941.0 4,449.3 (H I , He I , CO . µ m band, a few neutral metals) disappearedabout a year later.In the X-ray regime, the initial Jan-Mar 2005 data were pub-lished by Audard et al. (2005b). In brief, the X-ray data of early2005 indicated that the X-ray flux and luminosity stayed simi-lar within a factor of two during the outburst, and at the samelevel as in a pre-outburst observation in 2002. The fluxes inthe optical and near-infrared varied more significantly, withinfactors of − . The hydrogen column density showed noevidence for variation from its modest pre-outburst value of N H ≈ × cm − . However, there was evidence of a spec-tral change from a dominant hot plasma ( ≈ MK) in 2002and in January 2005 to a cooler plasma ( ≈ MK) in February2005 and probably in March 2005. We hypothesized that the hotmagnetic loops high in the corona were disrupted by the clos-ing in of the accretion disk due to the increased accretion rateduring the outburst, whereas the lower cooler loops were proba-bly less affected and became the dominant coronal component(Audard et al. 2005b). We argued that the cool component inV1118 Ori could not originate from shocks because free-fall ve-locities of matter falling from the truncation radius are too low(see Audard et al. 2005b for further discussion). In a subsequentpaper, Lorenzetti et al. (2006) independently analyzed our pub-lic data sets of early 2005, and also included the September 2005
XMM-Newton observation, together with their near-infrared datasets. The September 2005 observation showed a decrease in X-ray flux at the start of the decay phase.The initial paper by Audard et al. (2005b) presented and ana-lyzed the outburst data through March 2005. We present here the
XMM-Newton and
Chandra data obtained from September 2005on, which are analyzed in the context of our multiwavelengthanalysis. We also include a post-outburst data set obtained by
Chandra in December 2007. The latter was taken during a mi-nor outburst detected in the optical (Garcia & Parsamian 2008).Extensive optical and near-infrared photometry are presentedand analyzed, obtained from our team’s imaging data, along withmid-infrared photometry and spectroscopy with
Spitzer .
3. Observations and data reduction
Table 1 provides the observation log of our 2005–2006 monitor-ing campaign of the outburst of V1118 Ori with
XMM-Newton (Jansen et al. 2001) and
Chandra (Weisskopf et al. 1996). We also provide the information about the 2002 serendipitous obser-vation of V1118 Ori (an observation in 2001 with
XMM-Newton was reported in Audard et al. 2005b but is not mentioned heresince the star was not detected). A deep
XMM-Newton obser-vation was obtained in early 2006 that complements the shortmonitoring observations. Finally, we include the post-outburst
Chandra observation in December 2007 as well. Note that theangular resolutions of both X-ray satellites were not high enoughto separate the V1118 Ori binary.The
XMM-Newton data were processed with SAS 7.0.Standard procedures were applied. We used an extraction cir-cle of radius 20 ′′ for the source and a nearby background cir-cular region of 60 ′′ radius (40 ′′ for Sep 2005 and 35 ′′ for Mar2006). Event patterns lower than 4 and 12 were used only forthe European Photon Imaging Cameras (EPIC) pn and MOS(Str¨uder et al. 2001; Turner et al. 2001), respectively. The back-ground flux levels were high during all XMM-Newton observa-tions, and in particular in March and September 2005. As de-scribed in Audard et al. (2005b), we used only the MOS1 andMOS2 data for the March 2005 observation. The September2005 observation was so affected by the background that we lostabout half the exposure time in the EPIC pn, while there wasno MOS data available. The deep March 2006 observation wasaffected by a system failure at the Mission Operations Centreand the EPIC pn experienced full scientific buffer in the last partof the observation, explaining the reduced exposure time in theEPIC pn (78ks) compared to the EPIC MOS (90 ks). No OpticalMonitor data were taken with the
XMM-Newton
X-ray observa-tion due to the presence of the nearby bright Trapezium stars.The
Chandra data were processed with CIAO 3.3 andCALDB 3.2.3 . The task psextract was used to extract the spec-tra for V1118 Ori and the nearby background. We used a circleof 2 ′′ radius ( ≈ pixels) for the star and an annulus centeredat the position of the star but with radii of 10 and 60 pixels forthe background; our background area was, therefore, 212 timeslarger than the source extraction area. Our extraction radius forthe star includes 95% of the encircled energy at 1.5 keV and90% at 4.5 keV. Note that for the January 2005 observation insub-array mode, we used an outer radius of 40 pixels for thebackground annulus. For the 2002 observation, we used two cir-cles of radii of 15 ′′ and 45 ′′ for the source and the background,respectively (see Audard et al. 2005b). We used the pipeline data for the December 2007 observation,which was calibrated with CALDB 3.4.2. M. Audard et al.: The 2005 Outburst of V1118 Ori
Table 2.
Spitzer
IRAC and MIPS flux density measurements (units of mJy).
Date . µ m . µ m . µ m . µ m µ m2004 Mar 09 . . . . . . . . . . . . . . . . . ± . . ± . . ± . . ± . · · · · · · · · · · · · · · · . ± . . ± . . ± . . ± . . ± . · · · · · · . ± . · · · . ± . · · · . ± . . ± . . ± . . ± . · · · . ± . . ± . . ± . . ± . · · · . ± . . ± . . ± . . ± . · · · We observed V1118 Ori with the ANDICAM dual-channel im-ager on the SMARTS/CTIO Photometric coverage in the optical (standard Bessel
V RI ) wasobtained with the Celestron 14” optical tube assembly witha Paramount ME German equatorial mount at the VillanovaUniversity Observatory located on the campus near Philadelphia,PA. Observations were carried out with a SBIG ST7-XME de-tector thermoelectrically cooled to -25 degrees Celsius. Dark andflat field frames were collected at the end of each nights obser-vations. A standard error of 0.05 mag was estimated from thesignal-to-noise ratio and seeing conditions. Table A.3 (availableonline only) lists the magnitudes obtained at Villanova.
We have used published optical and near-infrared photometricdata from Lorenzetti et al. (2007) (
IJHK ), Garcia et al. (2006)and Garcia & Parsamian (2008) ( V ). In addition to our observations of V1118 Ori in outburst (pro-gram ID 3716, PI: G. Stringfellow) and in post-outburst (pro-gram ID 41019, PI: M. Audard), V1118 Ori was serendipitouslyobserved with the InfraRed Array Camera (IRAC; Fazio et al.2004) by the
Spitzer
Space Telescope (Werner et al. 2004) inMarch 2004 and twice in October 2004 (program IDs 43 and50, PI: G. Fazio). We show the IRAC images taken before theoutburst in March 2004 centered on V1118 Ori (Fig. 1). We SMARTS, the Small and Medium Aperture Research TelescopeFacility, is a consortium of universities and research institutions thatoperate the small telescopes at Cerro Tololo under contract with AURA. provide the IRAC fluxes for all observations in Table 2 (seealso Lorenzetti et al. 2007 for the IRAC data of October 27,2004). Details about the
Spitzer
IRAC data reduction are givenin Appendix B.V1118 Ori was also observed with the Multiband ImagingPhotometer for
Spitzer (MIPS; Rieke et al. 2004) at 24 µ mbefore the outburst, on March 20, 2004 (program ID 58, PI:G. Rieke). Again, Appendix B provides the details of the datareduction. No MIPS photometry is available for V1118 Ori dur-ing or after the outburst. Our programs (3716 and 41019) in-cluded MIPS spectral energy distribution data ( R ≈ ) in the70 µ m band; however, the on-time exposure (180 s) did not al-low us to detect V1118 Ori, even during the outburst. Indeed, thebackground level (due to diffuse emission in the Orion nebula)produced a much higher flux (of order Jy at 70 µ m) than theexpected signal from V1118 Ori (of the order 0.1 Jy if extrapo-lating the SED at 70 µ m, see Fig. 12). Spitzer also observed V1118 Ori with the InfraRedSpectrograph (IRS; Houck et al. 2004) twice during the outburst(PID 3716) with the Short-Low (SL: 5.2-14.7 µ m, R = λ/ ∆ λ ≈ ), Short-High (SH: 9.9-19.6 µ m, R = λ/ ∆ λ ≈ ), andLong-High (LH: 18.7-37.2 µ m, R = λ/ ∆ λ ≈ ) modules IRAC1
NEV1118 Ori
IRAC2IRAC3 IRAC4
Fig. 1.
Spitzer
IRAC images centered on V1118 Ori, taken pre-outburst in March 2004. The images are shown with a linearscale from 0 to 200 MJy sr − . The Herbig Ae star V372 Ori isseen near V1118 Ori. . Audard et al.: The 2005 Outburst of V1118 Ori 5
10 20 30 40Wavelength ( µ m)0.11.0 F ν ( Jy ) (d) 2008−11−14 F ν ( Jy ) (c) 2008−11−14 H − S ( ) H − S ( ) H − S ( ) [ N e II] [ S III] [ S III] [ S i II ] F ν ( Jy ) (a) 2005−02−18 SL SH LH F ν ( Jy ) (b) 2005−03−11 Fig. 2.
Spitzer
IRS background-subtracted spectra obtained near the peak of the outburst (panels a and b ) and after the outburst(panels c and d ). Panel d shows the SL and LL module data. Panels a , b , and c show the SL spectra as thick black lines, whereas theSH and LH spectra are shown as red and blue lines. Detected emission lines are labeled (while other lines well-subtracted by thebackground spectrum are shown in italics). Table 3.
Tentative
Spitzer
IRS line fluxes derived from the SH and LH spectra.
Line λ ( µ m) Flux ( − W m − )Feb 2005 Mar 2005 Nov 2008H − S (2) . . . . . . 12.2786 · · · · · · . ± . ii ] P / – P / . ± . . ± . . ± . iii ] P – P . . . . 18.7130 . ± . . ± . . ± . H − S (0) . . . . . . 28.2188 . ± . . ± . . ± . iii ] P – P . . . . 33.4810 . ± . . ± . . ± . (2005 February 18 and 2005 March 11). No background ob-servations were taken with the high-resolution modules. On theother hand, the post-outburst data (PID 41019; 2008 November14) were taken with background spectra for the high-resolutionmodules and also included the Long-Low (LL: 14.0-38.0 µ m, R = λ/ ∆ λ ≈ ) module.Figure 2 shows all IRS spectra after background subtrac-tion. We describe in Appendix B our methodology to derivethe background-subtracted spectra of V1118 Ori. We also dis-cuss in the Appendix the validity to use the post-outburst back-ground observation for our outburst IRS SH spectra. In brief,the post-outburst, background-subtracted high-resolution spec- trum is well-subtracted for the continuum, as the flux is con-sistent with the low-resolution spectrum. In the case of the out-burst SH spectra, the continuum is also accurate below 14 µ msince they are consistent with the low-resolution SL spectra.However, we prefer to stay on the safe side and claim that theobserved increase in continuum flux for λ > µ m duringthe outburst is unreliable, since we have no MIPS photometryor low-resolution spectra to confirm the increase. We also urgecaution with regards to the detection of lines in the background-subtracted SH and LH spectra. The strong background line emis-sion of the Orion nebula, and its inhomogeneity (see Fig. B.2)make the background subtraction difficult, although not impos- M. Audard et al.: The 2005 Outburst of V1118 Ori λ ( µ m)0.00.10.20.30.4 H − S ( ) [ N e II] F ν ( Jy ) λ ( µ m)0.00.10.20.30.40.5 H − S ( ) [ S III] λ ( µ m)0.00.10.20.30.40.50.6 H − S ( ) λ ( µ m)0.00.51.01.52.02.53.03.5 [ S III] [ S i II ] Fig. 3.
Zoomed in regions of the
Spitzer
IRS background-subtracted spectra. The post-outburst spectra are shown in the bottompanel, while the top two panels are for the outburst spectra. In the leftmost panels, we show both the high-resolution SH (thin,red) and the low-resolution SL (thick, black) data. Detected emission lines are labeled in the bottom panels (while other lines well-subtracted by the background spectrum are shown in italics). Observation dates applicable to each horizontal set of spectra are listedin the left panels.sible (e.g., a very strong PAH emission at 11.3 µ m is well-subtracted). The[S III ] emission lines at 18.71 and 33.48 µ m arelikely robust detections as their detection is consistent with thelow-resolution SL and LL data. The [Si II ] is also generally wellsubtracted. Interestingly, the rotational H S(0) line is detectedduring the outburst and in post-outburst. However the S(1) line isnot detected (it is is even slightly oversubtracted in the outburstspectra), and the S(2) line is only detected in the post-outburstspectrum (a faint excess in the outburst spectra is present only inone bin, likely due to slight differences in the wavelength scalesof the source and background spectra). It is likely that the H emitting conditions vary spatially in the Orion nebula, and thatthe background spectrum did not accurately reflect the H back-ground line fluxes near V1118 Ori, creating an excess in lineemission observed in the post-outburst spectrum. Alternatively,if the H emission lines are real, at least in the post-outburstspectra, the physical conditions near V1118 Ori are such thatthe para (odd quantum number J ) lines dominate the emission.A similar case of difficult background line subtraction arises inthe [Ne II ] line at 12.81 µ m: its detection in the SH spectrum isalso subject to caution, as it is not detected in the low-resolutionspectra (although it should have been, see the discussion in theAppendix). In any case, we provide in Table 3 the tentative, mea-sured line fluxes from the high-resolution module spectra, whileFigure 3 shows zoomed in regions of the IRS spectra of the de-tected lines.
4. X-ray analysis
Figure 4 presents the X-ray light curves for the
Chandra obser-vations. The lower panels show the binned light curves, whereas the upper panels show the photon CCD energy as a functionof time. Except for Sep 2002, the background is not subtractedin the light curves and not shown in the upper panels, since itcontributed to less than 1 count to the
Chandra total (source +background) number of counts over the total observing duration.Since V1118 Ori was observed serendipitously in the Sep 2002observation and was placed at the edge of the ACIS-I camera, thebackground contributed more significantly; therefore, the lightcurve shown in the bottom panel is background-subtracted, andthe CCD energy vs time background events are shown togetherwith the CCD energy vs. time total events (for clarity, everyninth background event was plotted, which corresponds approx-imately to the area-scaled background contribution).Small-scale variability is observed in the
Chandra lightcurves; however, no strong flares were detected during the ob-servations. The January 2006 observation displayed neverthelessa significant increase in the count rate in the last 10 ks, probablydue to a moderate flare. The average count rates are 3.14, 2.31,1.10, 0.39, 1.23, 1.23, and 1.09 ct ks − ( . − . keV, exceptfor Sep 2002: . − . keV) for the Sep 2002, Jan 2005, Jan-Jul2006, and Dec 2007 observations, respectively. While the effec-tive areas in all pointed observations are similar, the Sep 2002effective area was about 40-50% lower (due to the off-axis posi-tion of V1118 Ori). Therefore, to compare the above Sep 2002count rate with the later Chandra count rates, one needs to mul-tiply by a factor of about 2, i.e., 6.00 ct ks − . Clearly, assumingonly a change in emission measure, the X-ray flux dropped bya factor of 4–6, and even 15 in Feb 2006, during the 2005 out-burst compared to 2002. It is important to emphasize that theserendipitous observation did not show evidence of strong flar-ing, and V1118 Ori was likely caught in “quiescence”.Figure 5 presents the X-ray light curves for the XMM-Newton observations. In comparison with the
Chandra data, the . Audard et al.: The 2005 Outburst of V1118 Ori 7 E ( k e V ) µ ( c ks − ) Sep 2002 0.11.010.00 10 20 30 40 500.02.04.06.08.010.0 Jan 2005 0.11.010.0 E ( k e V ) µ ( c ks − ) Jan 2006 0.11.010.00 10 20 30 40 500.02.04.06.08.010.0 Feb 2006 0.11.010.0 E ( k e V ) µ ( c ks − ) Apr 2006 0.11.010.00 10 20 30 40 500.02.04.06.08.010.0 Jul 2006 0.11.010.0 E ( k e V ) µ ( c ks − ) Dec 2007
Fig. 4.
Light curves for the
Chandra observations (0.4–6.0 keVrange). The upper panels show the CCD energy of the events asa function of time after the start of the observations, whereas thelower panels show count rate light curves ( µ ) with a bin size of 1ksec. The average count rates are also shown as horizontal dottedlines. The time span for each panel was kept similar and equalto the time span of the Sep 2002 observation ( ≈ ksec). Sincethe other observations were shorter than 50 ksec, the rest of thetime span is marked as hashed regions. See text for a detaileddescription. XMM-Newton data were heavily impacted by background radia-tion. The (scaled) background contributions generally were sim-ilar or even higher than the net source contribution. In particular,during the March 2005 observation, the background completelyoverwhelmed the EPIC pn data which could not be used. Noobvious large flare was observed during the
XMM-Newton ob-servations, like in the
Chandra observations. However, in March2006, V1118 Ori showed a significant increase in X-ray flux inthe last 35–40 ks of the observation, probably due to flares or theonset of an active region. This increase is observed in all EPICcameras, comforting us that this behavior is not due to an in-crease in background flux that was improperly subtracted in theEPIC pn data.The light curves in Figures 4 and 5 suggest that the spectralfits represent snapshots of the thermal emission measure distri-bution of V1118 Ori’s corona observed serendipitously beforethe 2005 outburst, and over the course of the optical/infraredoutburst. µ ( c s − ) Feb 05 µ ( c s − ) Mar 05 µ ( c s − ) Sep 05 µ ( c s − ) Mar 06
Fig. 5.
Background-subtracted count rate light curves ( µ ) forthe XMM-Newton observations, with bin sizes of 500 s and inthe 0.4–6.0 keV range. The EPIC pn data are shown, except inMarch 2005 for which we summed the data of both MOS detec-tors. The (scaled) background light curves are shown as dashedcurves.
Figure 6 shows the long-term light curve of V1118 Ori in theX-rays, optical, and infrared, from its pre-outburst detection inSeptember 2002, through its outburst in 2005 and 2006 to thepost-outburst phase at the end of December 2007. We havelooked into correlations between the X-ray and optical and in-frared photometry. Since we obtained a few sequential opticaland infrared photometry points every night of our campaign, itwas not possible to determine short-term variations to comparewith the X-ray light curves. Nevertheless, to quantify further thecorrelation between the optical and near-infrared flux densitieswith the X-ray flux, we have calculated average optical and near-infrared flux densities within ± days of the 6 X-ray observa-tions between January 2005 and February 2006 (we did not in-clude the March 2006 observation, as this one might have beencontaminated by a flare). The optical and near-infrared flux den-sities are well-correlated with the X-ray fluxes measured at Earth(Fig. 7) as indicated by Pearson’s correlation coefficients, whichare in the range of ρ = 0 . − . . The sparsity of X-ray datapoints does not allow to measure any clear delay between the X-ray and optical/near-infrared flux densities. We have also donecross-correlation studies with the optical and near-infrared data(including only the data points for which we had simultaneousmeasurements). The onset of the outburst is unclear in most lightcurves (except perhaps in the V band), but the end of the outburstoccurs at a similar epoch for all bands (around MJD 53860).The shape of the light curves also differ: the fluxes vary most atshort wavelengths ( BV RI ), while the near-infrared fluxes dis-play shallower variations and they reach their peak some time
M. Audard et al.: The 2005 Outburst of V1118 Ori later (e.g., K peaks around MJD 53470 while V peaks aroundMJD 53433). The overall shapes of the light curves suggest thatthe mechanism dominating in the optical and near-infrared bandsmay be the same (e.g., due to a hot spot), although the differentpeak times suggest contamination of another mechanism long-ward of µ m, probably disk thermal emission (see below). Audard et al. (2005b) presented spectral fits of the V1118 Oridata taken through March 2005. Lorenzetti et al. (2006) repeatedthe analysis for the February 2005 data and added the September2005 data, noting a decrease in the X-ray flux of V1118 Ori atthat period compared to earlier in the outburst. For this paper,we have reprocessed the 2005 data with the latest calibration(see Sect. 3).We used XSPEC 11 (Arnaud 1996) to fit the
Chandra
ACIS and the
XMM-Newton
EPIC spectra for each epoch.For February 2005 and March 2006, we fitted the EPIC pn,MOS1, and MOS2 spectra simultaneously, while we used theEPIC MOS only for March 2005, and the EPIC pn only forSeptember 2005 (no MOS data were available). In general, weused a 1- T collisional ionization equilibrium (CIE) model ( apec ;Smith et al. 2001) with a photoelectric absorption model, exceptfor the deep March 2006 observation for which the signal-to-noise was large enough to use a 2- T CIE model. We have fittedthe plasma metallicity in the high signal-to-noise ratio spectraof February 2005 and used the best-fit value, Z = 0 . Z ⊙ , forthe other epochs. The coronal abundances are relative to the so-lar photospheric standard set of Grevesse & Sauval (1998). Thebest-fit metallicity is in line with the values measured for Fe inthe coronae of young stars in Orion (Table 3 in Maggio et al.2007). Since the spectral fits of the February and March 2005spectra are similar to those reported in Audard et al. (2005b),we provide in this paper their best-fit values (but adapt the emis-sion measure and luminosities to the adopted distance of 400 pc).Note that the September 2005 data are heavily contaminated bythe high background level during the observation. Contrary toLorenzetti et al. (2006), we preferred not to provide spectral fitsfor this observation; however, we provide estimates for the ob-served X-ray flux at Earth and absorption-corrected X-ray lu-minosity, based on the previous February 2005 and the posteriorMarch 2006 XMM-Newton observations. The results of our spec-tral fits are given in Table 4.
In view of the low observed count rates with
XMM-Newton and
Chandra during the outburst, we have explored an alter-native method to determine spectral properties for low-countX-ray spectra and have used the quantile analysis presentedby Hong et al. (2004) . This type of analysis was, in particu-lar, successfully used by Grosso et al. (2005) for their analysisof the spectral properties of V1647 Ori (see also Skinner et al.2009b). In brief, this method makes direct use of the event en-ergy values, and determines spectral properties based on quan-tiles of the total number of counts (e.g., median 50% quantile,and quartiles, 25% and 75%), instead of spectral energy bins.Such quantiles are then used as indicators for the X-ray color ofthe source. In particular, we make use of a diagram with the x =log ( m/ (1 − m )) index in the x -axis, where m corresponds to The code is available at .We used version 1.7. the median quantile Q , and the index y = 3 × Q /Q in the y -axis. Hong et al. (2004) define the α % quantile, Q α , as Q α = E α % − E lo E up − E lo , where E up and E lo are the lower and upper boundaries ofthe used energy band, and E α % is the energy below which thenet counts is α % of the total number of counts. Note that thismethod is able to take the distribution of background event en-ergies into account. We refer the reader to Hong et al. (2004) formore details.The total (source + background) and background extractionregions used in the quantile analysis were the same as those forspectral fits. The energy bands used to determine the extractionof events are reported in Table 4 together with the 25%, 50%, and75% quantiles, the E energy, and the x and y indices. We alsoprovide estimated values for the hydrogen column density andplasma temperature derived from the x and y indices. Finally,Figure 8 shows the color-color diagrams for all observations.There exists a certain degree of degeneracy in the x and y indices (e.g., Grosso et al. 2005): indeed, the same color in-dices can sometimes be described by different combinationsof column density N H and plasma temperature T . However,these degeneracies are limited to certain sets of T and N H : forexample, in the September 2002 observation, color indices of x = − . and y = 1 . can represent any plasma tempera-ture from about kT ≈ . to about kT ≈ . keV and N H can range from to . × cm − . At higher temperatures,the degeneracy breaks down, and the observed x/y values ofV1118 Ori can be better constrained to kT ≈ . − . keV and N H ≈ (1 − × cm − . The degeneracies differ slightlydepending on the response matrix generated in the observation.For example, with our Chandra monitoring data, the degenera-cies occur around ( x, y ) = ( − . , . and at ( − . , . .The first set again indicate a degeneracy for kT < . keV and N H < × cm − . The second set shows that the plasmaproperties cannot be disentangled between . and . keV for acolumn density of (3 − × cm − .In contrast to the temperatures obtained from spectral fits,the temperature derived from the color indices are more sta-ble and vary little from − MK, except in February 2006( T xy ≈ MK) and March 2006 ( T xy ≈ MK). In the caseof March 2006, the quantile temperature is in fact close to thesecond temperature component (which also has the higher emis-sion measure) obtained with the 2- T fit. In the case of February2006, the low number of counts did not allow a good constrainton the temperature. There also appears to have some kind ofdegeneracy: while the spectral fit gives a low plasma tempera-ture and “high” column density, the quantile analysis gives theinverse. It remains unclear which result is more robust, but thequantile diagram does not show evidence of a degeneracy. It isthen likely that the spectral fit reached a local minimum at a lowtemperature and could not find a combination of high plasmatemperature and low N H . The February 2005 quantile analysisand spectral fit both confirm that the temperature of the coro-nal plasma was cooler than in September 2002 and March 2005,as initially report in Audard et al. (2005b). The error bars of thecolor indices in January 2005 do not allow us to constrain thetemperature of the coronal plasma, while the spectral fit withXSPEC indicates no evidence of cool plasma at the start of theoutburst. . Audard et al.: The 2005 Outburst of V1118 Ori 9 −10 −5 l og F l u x ( a r b i t r a r y un i t s ) U IRVBJHKX−rayI1I2I3I4I1I2I3I4I1I2I3I4I1I2I3I4I1I2I3I4I1I2I3I4
Fig. 6.
Optical (
U BV RI ), near-infrared (
JHK ), mid-infrared (IRAC bands at 3.6, 4.5, 5.8, 8.0 µ m), and X-ray light curves ver-sus Julian date. Arbitrary scales were applied to help visualize the light curves. See text for the definition of the time intervalslabeled at the bottom. Data from the SMARTS and Villanova campaigns, from Lorenzetti et al. (2007), Garcia et al. (2006) andGarcia & Parsamian (2008) are included. We have attempted to search for correlations between X-rayproperties ( N H , T , N H ,xy , T xy , EM , L X , E ), but found gen-erally little or no correlation based on the Pearson correlationcoefficient, ρ . Indeed, its absolute value amounted to values typ-ically below . . One exception was a mild correlation betweenthe average median photon energy, E , and the absorbing col-umn density derived from the quantile analysis, N H ,xy , whichshowed ρ = 0 . . This is expected, as larger column densitieswill absorb low-energy photons. However, we emphasize that wehave only used the estimated N H ,xy based on the location of the xy values in the quantile color-color diagrams, and that we havenot taken into account the (relatively) large range of column den-sities covered by the quantile errors (typically, − × cm − ,similar to values derived from spectral fitting). Another excep- tion was a strong correlation between the emission measure andthe X-ray luminosity, with ρ = 0 . . Together with the abovelack of correlation between the plasma temperature and L X , and N H and L X , this is a strong indication that the X-ray variabil-ity observed during the outburst of V1118 Ori is only due to theamount of material in the corona. It is, thus, probably related tothe amount of mass falling from the disk into the stellar magne-tosphere, i.e., to the mass accretion rate.
5. Optical and infrared analysis
Figure 9 shows the silicate feature obtained at the three differentepochs and its normalized version ( S ν = 1 + F ν /F C ). We haveuse a linear fit to the SL spectra in the log λ − log F ν plane, using F l u x den s i t y ( Jy ) B V R I X (10 −14 erg s −1 cm −2 )10100 F l u x den s i t y ( Jy ) J H K Fig. 7.
Optical and near-infrared flux densities (averaged within ± days of the X-ray observation) as a function of the observedX-ray flux during the outburst (Jan 2005 to Feb 2006). Notice the different flux axis scales for the optical and near-infrared bands.only wavelength ranges from 6.0–8.0 µ m and 13.0–13.5 µ m todetermine the underlying continuum flux level, F C . The shapesof the silicate feature are compatible for all three observations,but the Mar 2005 silicate feature appears to be brighter than theFeb 2005 or post-outburst (Nov 2008) normalized fluxes, i.e.,more flux coming from the less optically thick disk upper layers,likely due to hotter temperature of such layers. The heating of theupper layers is probably due to the strong irradiation, e.g., by ahot spot (see below). The ratio S . /S . is equal to . − . ,suggesting grain growth, and the ratio ( S . /S . ) / ( S µm peak ) is . − . , evidence of a flat silicate feature.In the outburst spectrum of the EXor prototype EX Lup,´Abrah´am et al. (2009) observed a silicate feature with a crys-talline peak, in contrast to the pre-outburst silicate feature whichwas similar to the interstellar medium spectrum dominated byamorphous grains. They concluded that thermal annealing in thesurface layer of the proto-planetary disk was the mechanism forthis change in spectral features. Unfortunately, we have no pre-outburst spectrum of V1118 Ori before its outburst to comparewith the outburst spectra, only a spectrum taken 2 years after theend of the outburst and 3.5 years after our last outburst spec-trum in March 2005. If the pre-outburst spectrum was indeedtypical of amorphous silicates, it means that the time scale fordisappearance of the crystalline features must be longer than 2years. ´Abrah´am et al. (2009) proposed that similar crystalliza-tion must have occurred during the previous 1955-1956 outburstof EX Lup, and since the 2005 quiescent spectrum did not showcrystalline features, the removal time scale must have been lessthan 50 years. On the other hand, in V1118 Ori, it is also pos-sible that the unchanging silicate feature in the three spectra istypical of the quiescent silicate feature in the young star, per-haps an indication that the disk is more evolved. In any case,M-type young stars typically show flat silicate features (e.g.,Kessler-Silacci et al. 2005, 2006). Finally, we note that the IRSdata do not show strong evidence of the forsterite crystalline fea- ture at 16 µ m, in contrast with the feature observed in EX Lup( ´Abrah´am et al. 2009). This is consistent with the lack of a sharppeak at 10 µ m and of a shoulder at 11.3 µ m. We have calculated color indices from our optical and near-infrared data, and looked into their evolution during the out-burst. Time intervals were defined in the pre-outburst epoch( ; JD < ), initial phase of rise ( ; < JD < ), second phase of rise ( ; < JD < ),plateau ( ; < JD < ), decay ( ; < JD < ), and post-outburst ( ; < JD < ).Figure 10 shows the color-magnitude diagrams, whereasFigure 11 shows the color-color diagrams, together with an A V = 1 mag reddening vector using the reddening law ofRieke et al. (1985), the loci of the unreddened main sequenceand giant stars (Bessell & Brett 1988), and the loci of classicalT Tauri stars (Meyer et al. 1997). Compared to the pre-outburstcolor indices of V1118 Ori, the optical color indices are bluer inthe early phases of the outburst (except for B − V which shows apeculiar pattern, probably because the B magnitude is very sen-sitive to the evolution of hotspot emission), and gradually returnto the pre-outburst values in the decaying phase of the outburst.In the near infrared, an apparent blueing occurs as well in the ini-tial phases of the outburst with a return to “normal” colors in thedecay phase. The above results are consistent with the hypoth-esis of Lorenzetti et al. (2007) that blueing must have occurreddue to a hotter temperature component during the initial phaseof the outburst (their data were taken only in the decaying phaseand showed reddening). The variations in the color-color dia-grams also show no trend related to extinction variations, as sug-gested by Lorenzetti et al. (2007) and Lorenzetti et al. (2009).There is, however, a larger scatter in the infrared color-color dia-grams, especially in J − H versus H − K . Since the stellar pho- . Audard et al.: The 2005 Outburst of V1118 Ori 11 Table 4.
Spectral fits and quantile properties.
Parameter Sep 2002 Jan 2005 Feb 2005 Mar 2005 Sep 2005Satellite
Chandra Chandra XMM-Newton XMM-Newton XMM-Newton
Spectral analysis:Net counts a . . . . . . . . . . . . . . . . . . . . . . . . / . / . . / . / . . Scaled background counts a . . . . . . . . . . / . / . . / . / . . Exposure (ks) a . . . . . . . . . . . . . . . . . . .
85 4 .
66 18 . / . / .
91 (7 . / . / .
22 (13 . N H ( cm − ) . . . . . . . . . . . . . . . . +1 . − . . +3 . − . . +1 . − . . +0 . − . · · · T (MK) . . . . . . . . . . . . . . . . . . . . . . . . . +6 . − . + ∞− . +1 . − . . +3 . − . · · · EM ( cm − ). . . . . . . . . . . . . . . . . +0 . − . . +0 . − . . +1 . − . . +0 . − . · · · Z/Z ⊙ . . . . . . . . . . . . . . . . . . . . . . . . . . := 0 .
17 := 0 .
17 0 . +0 . − . := 0 . · · · F X ( − ergs cm − s − ). . . . . . . +0 . − . . +1 . − . . +0 . − . . +0 . − . . +1 . − . b L X ( ergs s − ) . . . . . . . . . . . . . . +0 . − . . +0 . − . . +0 . − . . +0 . − . . − . b Quantile analysis:Energy range . . . . . . . . . . . . . . . . . . . . . − . . − . . − . . − . . − . Net counts c . . . . . . . . . . . . . . . . . . . . . .
44 10 .
79 286 .
45 196 .
34 9 . Scaled background counts c . . . . . . . .
56 0 .
21 222 .
55 425 .
66 228 . Q . . . . . . . . . . . . . . . . . . . . . . . . . . . ± .
01 0 . ± .
07 0 . ± .
01 0 . ± .
02 0 . ± . Q . . . . . . . . . . . . . . . . . . . . . . . . . . . ± .
01 0 . ± .
07 0 . ± .
01 0 . ± .
02 0 . ± . Q . . . . . . . . . . . . . . . . . . . . . . . . . . . ± .
03 0 . ± .
14 0 . ± .
01 0 . ± .
03 0 . ± . E (keV) . . . . . . . . . . . . . . . . . . . . . ± .
06 1 . ± .
25 0 . ± .
04 1 . ± .
07 1 . ± . x = log[ Q / (1 − Q )] . . . . − . +0 . − . − . +0 . − . − . +0 . − . − . +0 . − . − . + ∞−∞ y = 3 × Q /Q . . . . . . . . . . . +0 . − . . +0 . − . . +0 . − . . +0 . − . . +1 . − . T xy (MK) . . . . . . . . . . . . . . . . . . . . . . ≈ ≈ ≈ ≈ · · · N H ,xy ( cm − ) . . . . . . . . . . . . ≈ ≈ ≈ ≈ . · · · Jan 2006 Feb 2006 Mar 2006 Apr 2006 Jul 2006 Dec 2007Satellite
Chandra Chandra XMM-Newton Chandra Chandra Chandra
Spectral analysis:Net counts a . . . . . . . . . . . . . . . . . . . . . . . . / . / . . . . Scaled background counts a . . . . . . . . . . / . / . . . . Exposure (ks) a . . . . . . . . . . . . . . . . . . .
82 27 .
11 78 . / . / .
95 29 .
57 29 .
57 34 . N H ( cm − ) . . . . . . . . . . . . . . . . +2 . − . . +4 . − . . +0 . − . . +5 . − . . +2 . − . . +1 . − . T (MK) . . . . . . . . . . . . . . . . . . . . . . . . . + ∞− . . +9 . − . . +3 . − . / +32 − . +8 . − . . +2 . − . . +7 . − . EM ( cm − ). . . . . . . . . . . . . . . . . +0 . − . . +0 . − . . +0 . − . / . +0 . − . . +0 . − . . +1 . − . . +0 . − . Z/Z ⊙ . . . . . . . . . . . . . . . . . . . . . . . . . . := 0 .
17 := 0 .
17 := 0 .
17 := 0 .
17 := 0 .
17 := 0 . F X ( − ergs cm − s − ). . . . . . . +0 . − . . +0 . − . . +0 . − . . +0 . − . . +0 . − . . +0 . − . L X ( ergs s − ) . . . . . . . . . . . . . . +0 . − . . +0 . − . . +0 . − . . +0 . − . . +0 . − . . +0 . − . Quantile analysis:Energy range . . . . . . . . . . . . . . . . . . . . . − . . − . . − . . − . . − . . − . Net counts c . . . . . . . . . . . . . . . . . . . . . .
48 10 .
53 626 . .
37 36 .
32 37 . Scaled background counts c . . . . . . . .
52 0 .
47 985 . .
63 0 .
68 1 . Q . . . . . . . . . . . . . . . . . . . . . . . . . . . ± .
04 0 . ± .
03 0 . ± .
010 0 . ± .
02 0 . ± .
01 0 . ± . Q . . . . . . . . . . . . . . . . . . . . . . . . . . . ± .
04 0 . ± .
09 0 . ± .
006 0 . ± .
03 0 . ± .
04 0 . ± . Q . . . . . . . . . . . . . . . . . . . . . . . . . . . ± .
08 0 . ± .
19 0 . ± .
019 0 . ± .
06 0 . ± .
03 0 . ± . E (keV) . . . . . . . . . . . . . . . . . . . . . ± .
15 1 . ± .
35 1 . ± .
06 1 . ± .
09 1 . ± .
14 1 . ± . x = log[ Q / (1 − Q )] . . . . − . +0 . − . − . +0 . − . − . +0 . − . − . +0 . − . − . +0 . − . − . +0 . − . y = 3 × Q /Q . . . . . . . . . . . +0 . − . . +0 . − . . +0 . − . . +0 . − . . +0 . − . . +0 . − . T xy (MK) . . . . . . . . . . . . . . . . . . . . . . ≈ ≈ ≈ ≈ ≈ ≈ N H ,xy ( cm − ) . . . . . . . . . . . . ≈ ≈ . ≈ . ≈ . ≈ . ≈ . The uncertainties are based on 68% Bayesian confidence ranges. X-ray luminosity and absorbed X-ray flux in the . − keV range, assuming d = 400 pc. a For
XMM-Newton observations, values refer to the EPIC pn, MOS1, and MOS2, respectively (pn only for September 2005). The values in parentheses refer todetectors not used in the spectral fits. b Estimates based on models for February 2005 and March 2006. c Based on pn data in general, except for March 2005 for which MOS1 data only were used (similar values and results were obtained for MOS2 data). tospheric emission does not vary during the outburst, we need tounderstand the variations in color indices and magnitudes dur-ing the outburst in the context of a star + disk + hot spot model.Lorenzetti et al. (2007) suggested that the polarization found inthe I band of the V1118 Ori outburst was intrinsic and that it wasprobably due to the spotted and magnetized stellar photosphere.As we will show below, the photometric changes can indeed beexplained by changes in the fluxes of the disk thermal emission and from a hot spot as the amount of mass falling from the diskonto the stellar photosphere increases during the outburst. Wenote also that during the outburst, any increase in hotspot emis-sion will further increase the irradiation of the disk. The follow-ing section aims to understand the multi-wavelength photometryand spectroscopy observed during the outburst of V1118 Ori bycomparing the data with young star disk models. Fig. 8.
Color-color diagrams based on the Q ( = m ), Q , and Q quantiles. The grid is in plasma temperature (kT in keV)and hydrogen column density ( N H in cm − ). The serendipitous, pre-outburst observation of September 2002 is shown at thetop, the post-outburst version of December 2007 is shown at the bottom. The 9 outburst observations in between are ordered in timefrom left to right and top to bottom, with January 2005 at the top left and July 2006 at the bottom right. Figure 12 shows the observed SED for V1118 Ori. We haveshown as small colored dots the data for the different phases(as defined at the beginning of section 5.2 and in the caption of Fig. 10), and the pre- and post-outburst photometry as redcrosses and black stars, respectively. The IRS low-resolutionspectra are also shown for the outburst phases (violet) and post-outburst phase (black). For clarity, we do not show the high-resolution spectra. To emphasize the expected stellar photo- . Audard et al.: The 2005 Outburst of V1118 Ori 13 λ ( µ m)0.100.150.20 F l u x ( Jy ) + ( F ν − F c ) / F c λ ( µ m) Fig. 9.
Spitzer
IRS SL spectra of V1118 Ori for the three epochs ( left panel). The dotted line shows the fit to the underlyingcontinuum in the [6-8,13-14.5] µ m range. The right panels show the normalized flux at the three epochs. B V R B I K K K V Fig. 10.
Magnitude versus color indices for the different time intervals (red: (e.g., Hauschildt et al. 1999) fora star with T eff = 3600 K, log g = 4 . , and Z = 0 (STARdusty2000 models), and reddened it with A V = 1 . magand the reddening law of Fitzpatrick (1999). Indeed, the hydro-gen column density derived from X-rays is . +1 . − . × cm − in 2002, and does not change much during the outburst, gener- ally staying in the range − × cm − , which correspondsto A V in the range A V = 1 . +0 . − . mag for R V = 3 . (Galacticvalue) or A V = 1 . +0 . − . mag for R V = 5 . (dusty environ-ment).We calculated the median magnitude differences from 0.44to 8.0 µ m (we have no U photometry during the outburst), rela-tive to the median magnitude in quiescence (we used both pre-outburst and post-outburst magnitude; Fig. 13). There is clear B − V V − R B − V B − V J − H J − K Fig. 11.
Color index diagrams for the same time intervals as in Fig. 10. An A V = 1 mag vector is added using the reddening lawof Rieke et al. (1985). The loci of unreddened main sequence and giant stars (Bessell & Brett 1988) are shown for the near-infraredpanels as solid and dotted curves, respectively. Finally the loci of classical T Tauri stars (Meyer et al. 1997) are shown as a dashedline in the ( J − H ) vs. ( H − K ) diagram.evidence that the optical emission brightened much more thanthe infrared emission, with a peak in the V band. We also ob-serve that the BV RI decayed faster than the
JHK bands, prob-ably indicating a mixture of different emission components inthe near-infrared bands (e.g., disk and hot spot). At long wave-lengths ( λ ≥ µ m), the increase in flux was similar at all wave-lengths.To quantify better the star+disk properties of V1118 Ori be-fore and during its outburst, we have then used the code de-scribed in Whitney et al. (2003a,b), version 20090224 . We haverun the code by fixing certain parameters and leaving othersfree to vary. First, the models did not include an envelope, butonly an ambient medium with density of − g cm − . Thecontribution of this medium is generally negligible. The exclu-sion of the envelope is motivated by the fact that it would con-tribute mostly at long wavelengths where we have no data. TheIRS data also do not show strong evidence of an envelope (thatcould not be attributed to improper background subtraction). Thefixed parameters were the stellar radius ( R ⋆ = 1 . R ⊙ ), mass( M ⋆ = 0 . M ⊙ ), and effective temperature ( T eff = 3600 K; see §
2; we used as input a lower spectral resolution version of thesame PHOENIX stellar atmosphere model as above), the disk http://gemelli.colorado.edu/˜bwhitney/codes/ mass ( M Jup ) and outer radius (
AU), the disk flaring power( ∝ R β , β = 1 . ) and the disk scale height ( . R ⋆ at R ⋆ ; theheight is thus . AU at 100 AU), and the disk density profile( ∝ R − α , α = 2 . ). For the dust opacity, we used a model ofgrains as large as 1 mm for the denser regions of the disk (seeWood et al. 2002 for HH 30), but used an interstellar mediumgrain model for the less dense regions of the disk (i.e., below athreshold set at . × − cm − ). The disk mass and outerradius were not constrained from our optical and infrared data .Changes in their values would also not affect much the mod-els at optical-infrared wavelengths. We have left the mass accre-tion rate and inner disk radius (in units of the dust destructionradius, R sub ) free to vary. The model also includes the possi-bility for material to fall from the disk onto a hot spot by fol-lowing the magnetic field lines. We have kept the truncation ra-dius to AU, but left the fractional area, f S (ratio of the shock We have obtained data with the IRAM 30m MAMBO2 bolometer at1.2 mm in late November 2009. V1118 Ori is detected at . ± κ ν = 0 . cm g − , us-ing a gas-to-dust ratio of 100, we obtain a disk mass of . ± . M Jup for a characteristic temperature of 20 K. The disk mass increases to . ± . M Jup for T c = 10 K. Although the determination of the exactdisk mass is out of scope in this paper, the derived estimates indicatethat the assumed disk mass in our model is realistic.. Audard et al.: The 2005 Outburst of V1118 Ori 15 µ m)0.0010.0100.1001.00010.000 F l u x ν F ν ( − e r g s − c m − ) Fig. 12.
Compilation of the optical and infrared photometry taken before, during, and just after the outburst (see text and caption ofFig. 10 for details about colors). The IRS SL and LL spectra and a reddened ( A V = 1 . mag) stellar atmosphere model are alsoshown.area to the star area; see Calvet & Gullbring 1998), free to vary.Finally, the code samples the model at 10 different inclinationangles from cos i = 0 . to . every . , i.e., from ◦ to ◦ . Thus, we have four different parameters that can vary( [cos i, ˙ M , f S , R min ]).Note that the choice of α and β is typical of those observed inT Tauri stars, and correspond to values that seem to fit the V1118Ori data. The power law indices are related in the sense that ifthe disk surface density is Σ ∝ R p , then p = β − α . A valueof β = 1 . is close to the theoretical values for irradiated and α disks ( β = 9 / . ), and, with α = 2 . , p = β − α = − . ,i.e., typical for a steady optically thick disk. Note that a fit tothe IRAC and MIPS fluxes in quiescence gives λF λ ∝ λ − . ,i.e., consistent with the commonly observed λF λ ∝ λ − / forirradiated disks. Assuming that the temperature follows a powerlaw T ∝ R − q , and if λF λ ∝ λ s , both indices are related by s = 2 /q − . Thus, we have q = 0 . , i.e., T ∝ R − / .There is no pre-outburst IRS spectrum available to determinethe shape of the silicate feature before the outburst (emission orabsorption). However, we have used the post-outburst IRS spec-tra, scaled by a factor of 0.7 to match the IRAC photometry (seeFig. 9). For the photometry, we have used the median of the pre-outburst (interval µ m flux without correction and shows thesilicate feature in emission.We ran the code with 5,000,000 “photons” or energy pack-ets for each model. This results in SED models with reasonablygood signal-to-noise ratios over the wavelength range covered byour data. The models were then reddened with A V = 1 . mag(assuming ISM opacities). For each realization we have calcu-lated a reduced χ p value χ p = 1 N N X i =1 (cid:18) log F ν ( λ i ) − log M ν ( λ i , p )∆(log F ν ( λ i )) (cid:19) (1)where F ν ( λ i ) are the data flux densities, M ν ( λ i , p ) are the model flux densities (for pa-rameter vector p = [cos i, ˙ M , f S , R min ]) in-terpolated at the wavelength of the data points( . , . , . , . , . , . , . , . , . , . , . , . µ mfor the photometry, and . , . , . , . . . , . µ m for the IRS– we only used the SL data). For the photometric flux densityuncertainties, we have used a fixed ∆(log F ν ( λ i )) = 0 . ,i.e., about 10% relative uncertainties for all data points (a valuelarger than many actual uncertainties on the measured values)to account for flux variations observed before and at the peak of µ m)0−1−2−3−4−5 M ed i an ∆ M ( m ag ) Fig. 13.
Median magnitude difference relative to quiescence as afunction of wavelength. The same time intervals as in Fig. 10 areused for the color coding (lime filled circles: V band. Fig. 14.
Probability distribution of the parameters( [cos i, ˙ M , f S , R min ]) for the quiescent (solid) and outburst(dotted) cases.the outburst. We used ∆(log F ν ( λ i )) = 0 . instead for theIRS SL data to put more weight on the silicate feature.Figure 14 shows the normalized model probabilities (calcu-lated as p p = e − χ p ) for the quiescent case (solid), and for theoutburst case (dotted). The distributions of inclination angles areslightly different for the quiescent and peak of outburst cases, butthe distributions overlap in the range cos i = 0 . − . , with amaximum probability (obtained by multiplying the probabilitydistributions) at cos i = 0 . , i.e., i ≈ ◦ . The mass accretion rate distributions are sharply peaked: the most probable valuesof ˙ M are − . M ⊙ yr − and − . M ⊙ yr − in quiescenceand and at the peak of the outburst, respectively, indicating anincrease by a factor of 4. The coverage factor seems to be small( f S < . ) outside the outburst while the simulations indi-cate high values ( . < f S < ) at the peak of the outburst. Notethat this parameter is essentially constrained by the optical andnear-infrared photometry. The minimum disk radius has likelyvalues of R min = (10 . − . ) R sub , with the most probablevalue at R min ≈ R sub , whereas the values at the peak of theoutburst are lower and most likely equal to the sublimation ra-dius. The best models give sublimation radii of about 0.10 AUand 0.20 AU in quiescence and at the peak of the outburst, im-plying inner disk radii of . − . AU (most probable valueof 0.4 AU) and 0.2 AU, respectively. This result suggests thatthe inner disk radius (determined from the infrared data arisingfrom dust emission) probably diminished during the outburst,although we emphasize that quiescent values of 0.2 AU are pos-sible, though less probable from our SED modeling.Figure 15 shows the best fits to the “quiescent” SED and tothe peak of outburst SED, respectively. The total (bolometric)luminosity (including all components) varied from about . L ⊙ in quiescence (with a contribution of the stellar luminosity, L =4 πR σT , of . L ⊙ and of the thermal disk luminosity of . L ⊙ ) to ≈ . L ⊙ at the peak of the outburst (the thermaldisk luminosity increasing to about . L ⊙ ).Thus the bolometric luminosity is dominated by the stellardirect emission, i.e., the stellar photospheric emission and thehotspot emission (the latter in particular during the outburst). Wenote, however, that we would need data at shorter wavelengthsthan B to constrain better the hotspot emission during the out-burst.
6. Discussion
Our multi-wavelength campaign allows us to obtain severalpieces of information on the outburst of V1118 Ori, as seen inthe optical and infrared, and in the X-rays. Overall, the evolu-tion of the X-ray flux during the outburst followed the evolu-tion of the optical and near-infrared fluxes. The sampling of theX-ray observation does not allow us to determine whether theincrease in X-ray flux in March 2006 is related to the short-term increase in optical/infrared flux observed about 45 daysearlier. The thermal properties of V1118 Ori showed evidenceof a plasma with temperature of a few MK to a few tens of MK.In our initial paper, we reported a change in the spectral proper-ties from a predominantly hot (20 MK) plasma before the out-burst (and possibly at the very beginning) to a lower temperature(8 MK) in February 2005 (Audard et al. 2005b). Our reanalysisconfirms this finding and a gradual return to larger values in thelater phases of the outburst. However, the low signal-to-noise ra-tio of the X-ray spectra taken during our monitoring campaignmakes it difficult to constrain accurately the plasma temperature(in contrast, the February and March 2005 observations had amuch better signal-to-noise ratio). The alternative quantile anal- Lorenzetti et al. (2006) quote a peak luminosity of . L ⊙ ; how-ever, they used the peak magnitudes of a previous outburst to build theoutburst peak SED and to derive the luminosity. The previous outbursthad a peak magnitude of J = 8 . , in contrast to the 2005–2006 outburstthat peaked at J = 10 . . In addition, our method derives the lumi-nosity by matching the observed SED with a detailed radiative transfermodel, while Lorenzetti et al. (2006) integrated the flux densities. Thetwo above issues may explain the difference in derived luminosities.. Audard et al.: The 2005 Outburst of V1118 Ori 17 Fig. 15.
The best-fit models to the “quiescent” SED (left) and to the peak outburst SED (right). The direct stellar contribution(including emission due to the hot spot) is shown as a dotted line (the A V = 1 . mag reddened stellar atmosphere model is shownas a dash-dotted line), the thermal disk spectrum as a dashed line.ysis showed evidence of less variations in the plasma tempera-ture. Nevertheless, it is consistent with the detected change intemperature. While V1647 Ori showed a strong increase in X-ray flux during the outburst (Kastner et al. 2004; Grosso et al.2005), V1118 Ori only showed moderate X-ray flux increase,and possibly an eventual decrease in flux after the outburst, ifwe assume that the 2002 flux level was representative of the pre-outburst X-ray flux. In this case, the impact of the mass infallonto the stellar magnetosphere must have significantly affectedthe latter, leaving it in a lower X-ray flux state. It is, however,also possible that the 2002 flux was observed by chance in ahigh quiescent state, and that the 2007 X-ray flux is closer tothe pre-outburst flux level. In any case, the changes in the X-rayfluxes during the outburst, and possibly the temperature changesall indicate a strong interplay between the material falling fromthe disk due to the increase in mass accretion rate and the stellarmagnetosphere.From the optical and infrared SED data and Monte-Carlomodels using the code of Whitney et al. (2003a,b), there is clearevidence that the thermal disk emission dominates in the mid-infrared in quiescence and during the outburst, while emissionfrom a hotspot and reprocessed emission dominates the outburstSED in the optical and near-infrared, and contributes to some ex-tent before the outburst. However, we note that this conclusiondepends on our choice of using disk models that stop at the sub-limation radius, i.e., there is no disk contribution at shorter radii,where the gas is optically thin. This could contribute to the near-infrared fluxes, but it should not contribute significantly to theoptical fluxes. Therefore, we believe that the addition of hotspotemission is necessary to fit the blue side of the V1118 Ori SED,during the outburst and probably before the outburst.The increase in mass accretion rate during the outburst leadsto the change in inner disk radius which moves closer to thestar’s surface. Such an effect could explain the changes in theX-ray emission (see above). From our simulations, we have de-termined that the mass accretion rate increased during the out-burst: consequently, the dust sublimation radius increased from ≈ R ⋆ = 0 . AU to ≈ R ⋆ = 0 . AU, whereas the innerdisk radius (based on the infrared dust emission) moved from . AU in quiescence to . AU in outburst (although the op-tically thin gas disk could be closer to the star), and the innerdisk radius is consistent with the dust sublimation radius at thepeak of the outburst. We emphasize that such estimates of the inner disk radius are biased due to our use of dust thermal emis-sion for the disk. The optically thin gas disk is difficult to de-tect and we would require gas tracers, e.g., H or CO bandheadsto properly determine the inner extent of the disk. In any case,the strong increase in fractional area of the hotspot suggests thatlarge amounts of matter fell from the disk onto the star during theoutburst, incidentally interacting with the stellar magnetosphere.Using their near-infrared spectrum taken in September 2005(i.e., about at the end of the plateau phase) and a wind lossmodel, Lorenzetti et al. (2006) derived a mass loss rate of about × − M ⊙ yr − from H I . Furthermore, little or no ab-sorption was reported ( A V < ), with no evidence of varia-tions during the outburst. They also derived a mass loss rateof (3 − × − M ⊙ yr − from the CO overtone emis-sion, indicating an ionization fraction of . − . . Furthermore,Lorenzetti et al. (2007) observed V1118 Ori in September 2006,i.e., well after the outburst, and detected no emission line, whilethe 2.3 µ m CO band was seen in absorption . Evidence of windemission in H α (P Cyg profile) was also reported by Herbig(2008) in November 2005, in contrast with the more symmet-ric profile observed after the outburst in December 2006. It is,therefore, likely that high wind loss rates occurred transientlydue to the instability in the disk. We note that disk emission mayprovide a better fit to CO bands in young stars than mass lossemission in general (e.g., Chandler et al. 1995), but this may notbe the case during the outburst of a young star such as V1118Ori: our SED models including emission due to matter fallingfrom the accretion disk onto the stellar photosphere clearly showthat the outburst SED at wavelengths below µ m is dominatedby the hotspot emission, not by disk thermal emission (while itremains negligible outside the outburst and stellar photosphericemission dominates). Wind loss emission may thus be present inthe optical and near-infrared.We can compare the mass loss rates derived for the neu-tral gas by Lorenzetti et al. (2006) to our derived mass accretionrates: the SED modeling suggests a variation in the mass accre-tion rate from about . × − M ⊙ yr − to . × − M ⊙ yr − ,i.e., these mass accretion rates are comparable to the mass loss Lorenzetti et al. (2006) derive the mass loss rate from the windmodel of Carr (1989) and his Figure 8. While their derived mass lossrate is accurate, they incorrectly give a CO luminosity of . × − L ⊙ ,while it should have been . × − L ⊙ based on the quoted line flux(D. Lorenzetti 2009, priv. comm.).8 M. Audard et al.: The 2005 Outburst of V1118 Ori rates determined by Lorenzetti et al. (2006). This is somewhatsurprising, since mass outflow rates are expected to be about 0.1times lower than mass accretion rates (Hartigan et al. 1995, withrevised mass accretion rates, see Gullbring et al. 1998; Edwards2009). Our mass accretion rate in quiescence is larger by abouta factor of 10 compared to those derived for classical T Tau stars( ≈ − M ⊙ yr − ; e.g., Gullbring et al. 1998), and lower by afactor of 10 than the rates derived for embedded Class I stars( ≈ × − M ⊙ yr − ; Kenyon et al. 1993). Therefore, our de-rived mass accretion rates are consistent with placing V1118 Oriin a category between Class I and CTTS stars. We note also thatthe IRAC color indices of V1118 Ori in quiescence are [3.6]–[4.5]=0.65–0.75 and [5.8]–[8.0]=0.65–0.7 which places it for-mally in the region of Class II sources, but also close to the re-gion for Class 0/I (e.g., Allen et al. 2004). Furthermore, the K –[3.6] color of 1.52 in “quiescence” also place the star as a strongaccretor. In contrast, U − V = − . which formally would(incorrectly) place V1118 Ori as a weak-lined T Tauri star. Infact, using the Gullbring et al. (1998) method to derive mass ac-cretion rates from the U -band luminosity excess, we find a massaccretion rate of . × − M ⊙ yr − , much smaller than val-ues derived by SED modeling, and also inconsistent with thestrong infrared excess. While our derived mass accretion ratesare model-dependent, we are confident that our derived mass ac-cretion rates are realistic, despite the difficulty to derive massaccretion rates from infrared excess (e.g., Kenyon & Hartmann1987).Overall, the data provide ample evidence that a significantportion of matter fell from the disk onto the young star due tothe increase in mass accretion rate. Over a duration of about400 days, we estimate about × − M ⊙ = 0 . M ⊕ was de-posited (using an average of half the peak mass accretion rate),i.e., about twice the amount that would have been deposited overthe same time range if using the quiescent mass accretion rateThe hotspot covering factor we derive at the peak of the outburstis also relatively large ( f S > . ), suggesting that a large frac-tion of the stellar magnetosphere may also have been influencedby the increase in mass accretion rate, which in turn supports thehypothesis that the stellar magnetosphere was modified duringthe outburst and the X-ray properties of the coronal plasma havechanged Audard et al. (2005b). Our SED modeling of the dustthermal emission suggests a change of the inner dust disk radiusfrom about 0.4 AU in quiescence to 0.2 AU at the peak of theoutburst (the optically thin gas could reach down to 0.2 AU orbelow before the outburst, although it would be unclear why dustwould not go down to the sublimation radius, unless, e.g., diskclearing by a giant planet or significant grain growth making thedust optically thin would be present in the inner portion of thedisk). Together with the increase in f S , this provides evidencethat the disk closed in during the outburst, disrupted the mag-netosphere, and matter fell onto the stellar photosphere, produc-ing strong hotspot emission that dominated the optical and near-infrared in outburst (while the thermal disk emission increasedby a factor of 4 and dominated the mid-infrared emission).While color-color and color-magnitude diagrams can be usedin young stars to determine the level of reddening, such diagramsmight be difficult to use to measure any reddening during theoutburst of a young star: in addition to reddening, the whole SEDof a young star changes dramatically, increasing in flux, but alsoin shape: in the case of V1118 Ori, the optical and near-infraredfluxes increased much more than the mid-infrared spectrum. Theformer are dominated by stellar and hotspot emission, while thelatter is dominated by thermal disk emission. Color-color dia-grams show a clear blueing of the colors (Fig. 11), not due to clearing of any enshrouding material, but due to the strong in-crease in flux below . µ m (Fig. 13). Such multi-wavelengthdata show the importance that a good coverage of the SED isnecessary to understand the disk-star interactions in young stars.Furthermore, Figures 10 and 11 indicate that V1118 Ori returned(within uncertainties) to its “initial”, pre-outburst conditions.
7. Conclusions
We have followed the young accreting star V1118 Ori before,during, and after its outburst in 2005–2006 with photometry in
U BV RIJHK and in X-rays with
XMM-Newton and
Chandra .We have further obtained
Spitzer photometry and spectroscopyduring the outburst and after the outburst.The optical and infrared data showed significant variations:the data below µ m brightened by as much as about − mag,while the infrared data brightened only up to ≈ mag. Monte-Carlo simulations of a star+disk+hotspot model suggested thatthe optical data were dominated by hotspot emission at the peakof the outburst (while the stellar photosphere dominated beforethe outburst), while thermal disk emission from optically thickdust dominated the infrared. The SED analysis showed an in-crease in mass accretion rate from ˙ M = 2 . × − M ⊙ yr − in quiescence to ˙ M = 1 . × − M ⊙ yr − at the peak of theoutburst, and an increase of the hotspot fractional coverage areafrom f S ≤ − to f S > . . The SED modeling suggests thatthe dust sublimation radius increased from . to . AU. Basedon the thermal dust emission, the inner disk radius moved fromabout 4 times the dust sublimation radius (0.4 AU) to 0.2 AU( = R sub ) at the peak of the outburst. But the gas inner disk radiuscould extend below the dust sublimation radius. The inclinationangle of the disk is constrained to values of cos i = 0 . − . with a peak at . , i.e., i ≈ ◦ .The Spitzer
IRS spectra showed the silicate feature at 10 µ min emission during and after the outburst, with no apparent vari-ation in shape or flux, indicating that the outburst had little im-pact on the optically thin region where the silicate emission isoriginating. Emission lines were detected in the high-resolutionmodule data although the background subtraction proved to bedifficult, especially for the outburst spectra. Follow-up observa-tions, e.g., from the ground should confirm the detection of linesof particular interest, e.g., [Ne II ], [S III ], and H − .The initial X-ray observations (Jan–Mar 05) indicated achange in the thermal structure of the coronal spectrum of V1118Ori with little change in the X-ray flux (Audard et al. 2005b),while fading in the X-rays was already detected in September2005 (Lorenzetti et al. 2006). Our continuing monitoring with XMM-Newton and
Chandra provided the complete picture, con-firmed the low plasma temperature measured in February 2005,and showed that V1118 Ori’s X-ray flux was strongly correlatedwith the optical and infrared fluxes, although intrinsic variabil-ity (due to magnetic activity) was also observed. Compared tothe 2002 flux, the increase in X-ray flux was not large duringthe outburst and the post-outburst X-ray observation in late 2007showed a lower flux level than observed serendipitously in 2002,which could either indicate that the outburst significantly im-pacted the stellar magnetosphere to leave it in a lower state thanbefore the outburst, or that the 2002 observation caught V1118Ori in a quiescent high state. The X-ray spectra showed evidenceof thermal variations of the corona during the outburst, but thesignal-to-noise ratio of our data was not large (except in March2006 during which we could detect a second temperature com-ponent, probably due to the increase in flux at the end of the . Audard et al.: The 2005 Outburst of V1118 Ori 19 observation, likely contaminated by a flare). After the apparentchange to a cooler plasma temperature in the early phase of theoutburst, we argue that the thermal properties of V1118 Ori’scorona returned to its “normal”, hotter values before the end ofthe outburst, but that the continuing mass infall affected signifi-cantly the magnetospheric configuration.Our observations have shown the interplay between disk andstellar magnetosphere in young accreting stars, as demonstratedby the changes in the X-ray plasma emission due to the strongincrease in mass infall from the disk onto the stellar photosphere.While our data do not directly help to determine the origin of theoutbursts (thermal disk instabilities or instabilities due to closecompanions or giant planets in the disk), it is interesting thatReipurth et al. (2007) identified V1118 Ori as a close visual bi-nary. Future observations should, ideally, take into account thisbinarity to determine which star is actually outbursting.
Acknowledgements.
We acknowledge support by NASA through
Chandra award DD5-6029X and through
XMM-Newton award NNG05GI96G toColumbia University, SAO GO6-7005D, SAO GO8-9020X, and NNX07AI37Gto the University of Colorado. M. A. acknowledges support from a SwissNational Science Foundation Professorship (PP002–110504). G. S. Stringfellowfurther acknowledged support from NASA programs NNX06AG44G,NNX07AT30G, and
Spitzer through JPL contract 1265288. The
Chandra
X-rayObservatory Center is operated by the Smithsonian Astrophysical Observatoryfor and on behalf of NASA under contract NAS8-03060. Based on observationsobtained with
XMM-Newton , an ESA science mission with instruments andcontributions directly funded by ESA Member States and NASA. This work isbased in part on observations made with the
Spitzer
Space Telescope, which isoperated by the Jet Propulsion Laboratory, California Institute of Technologyunder a contract with NASA. Support for this work was provided by NASA.Stony Brook’s participation in SMARTS is made possible by support from theoffices of the Provost and the Vice President for Research. We thank J. AllynSmith, P. McGehee, J. Espinoza, and D. Gonzalez for doing the observationswith the SMARTS telescopes. We also thank H. Tannanbaum and N. Schartelfor granting
Chandra
DDT and
XMM-Newton
TOO time to observe V1118 Ori.We also thank N. Grosso for discussions on quantiles and outbursting sources,A. Carmona and F. Fontani for useful comments on the manuscript. Finally,we thank the referee, D. Lorenzetti, for comments that clarified aspects of themanuscript.
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Appendix A: Details on the SMARTS datareduction
The ANDICAM optical detector is a Fairchild 447 2048 × µ m pixels. It is used in 2 × − and about a ′ × ′ field. Readnoise is 6.5 e − RMS; gain is 2.3 e − /DN. It is readwith a single amplifier. Filters are Johnson B and V and Kron-Cousins R and I . We generally took 3 exposures in B and V ,and 2 images in R C and I C , and co-added these prior to analysis.The IR channel detector is a Rockwell 1024 HgCdTe “Hawaii”array with 18 µ m pixels binned 2 × . ′′ pixel − plate scale. The IR channel field is . ′ × . ′ . The filters we usedare standard CIT/CTIO JHK . We flatten the images using domeflats obtained every 3 nights. We observe using a 3-point ditherpattern and a ′′ throw, and use the median of the unshiftedimages as the local sky. A.1. Differential Photometry
We initially determine relative magnitudes using differentialphotometry. In the optical, we selected the 11 brightest stars inthe ≈ ′ field of view (Table A.1; V372 Ori is overexposed).Instrumental magnitudes are determined for the target and allthe comparison stars, by summing the counts in a 5 pixel ( . ′′ )aperture. Background is the median level in an annulus from 10-20 pixel ( . ′′ − . ′′ ) radius. Since V1118 Ori is in a region ofcomplex nebulosity, we also tried modelling the background bytaking radial cuts through the image and extrapolating the levelat the position of the target. This affected the magnitudes at lessthan 5% at B , and even less at longer wavelengths.In the near-infrared, we use 6 comparison stars in the field(again excluding V372 Ori). We use a 5 pixel ( . ′′ ) aperture.Background is the median level in an annulus from 10-20 pixel( . ′′ − . ′′ ) radius.We construct a single comparison, the “superstar” from aweighted sum of all the comparison stars. The differential mag-nitude is the difference between the instrumental magnitudes ofthe target and the “superstar”.As it happens, most of the stars in this field are known or sus-pected variables. Fortunately, none are as variable as V1118 Ori.We assume that, by summing all the comparisons, the “super-star” will be less variable than any single comparison. We haveconstructed light curves using subsets of the comparison stars,and do not see any gross effects that can be attributed to varia-tions of the comparisons. Nonetheless we have not incorporatedthose stars with measured standard deviations after absolute cal-ibration (see below) > V into the optical “superstar”.We caution that there is no certainty that the comparison is stableat the few percent level. A.2. Optical Absolute Calibration
Although the SMARTS operations are not optimized for abso-lute photometry, single observations are made of Landolt stan-dard fields on each night judged photometric to establish thephotometric zero point. These can be used to determine apparentmagnitudes with an accuracy of better than 10%.On each night on which we have both observations of thetarget and of a photometric standard field, we compute the zero-point offset (apparent magnitude - instrumental magnitude). Weassume the standard atmospheric extinction correction for CerroTololo. Since the standard field is observed only once per night,we cannot account for changes in transparency through the night. We do not solve for reddening terms, so the photometric solutionconsists of only the zero-point.We use the zero-point correction to determine the apparentmagnitudes of all the comparison stars in the field. We could dothis on 71 nights. In doing so, we confirm that all the comparisonstars are variable (Table A.1). We compared these magnitudesto the
V, I C magnitudes published by Hillenbrand (1997). Themedian offsets are 0.08 and 0.23 mag at V and I C , respectively.Given that all the stars are variable members of the Orion Nebulapopulation, this is acceptable agreement.We apply this absolute calibration on all nights (includ-ing non-photometric nights) to determine the apparent magni-tudes of the target. There are 14 nights where we overlap withthe Villanova photometry. On these nights there are system-atic offsets, in the sense (SMARTS-Villanova), of 0.16 ± ± ± V , R C , and I C , respectively.The cause for the systematic offsets may be the lack of a colorterm in the photometric solution, exacerbated by the fact thatmost of our comparisons have V − I C colors > A.3. Near-infrared Absolute Calibration
For the
JHK absolute calibration, we compare the instrumentalmagnitudes directly to the 2MASS magnitudes of the 6 com-parison stars. To check for target variability, we have comparedthese magnitudes to those published in the DENIS catalog, andby Ali & Depoy (1995) and Hillenbrand et al. (1998). Two ofthe six comparisons (2MASS J05344159-0534249 and 2MASSJ05344219-0533036 appear variable, with variances > K (4 observations), and 2 others (JW 94 and JW 100) appearvariable in the optical (Table A.1).Our ability to absolutely calibrate the photometry is ulti-mately limited by the stability of the zero-point. We trust inregression to the mean, that by using at the possible compar-isons, on average the zero-point will be more stable that de-termined from any single comparison. We do not include thesystematic uncertainty in the zero-point in our error budgets.Overall, the JHK fluxes compare well to those published byLorenzetti et al. (2007). However, we have added a systematiccorrection of +0 . mag for J and +0 . mag for K to matchtheir photometry. V1118 Ori is intrinsically variable, as are thecomparison stars. The only reason for applying the systematicoffsets to our data are to simplify comparison graphically withthe Lorenzetti et al. data. We note that the offsets do not changesignificantly the results of this paper, they merely shift the near-infrared data points by . to . mag, e.g., in Figures 10 and11. . Audard et al.: The 2005 Outburst of V1118 Ori , Online Material p 2
Table A.1.
Comparison stars for the SMARTS photometry
Comparison
B V R I δV δI
IV Ori . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 16.19 ± ± ± ± · · · · · · CSV 6218 . . . . . . . . . . . . . . . . . . . . . . . . . . 17.40 ± ± ± ± ± ± ± ± ± ± ± ± · · · · · · V1175 Ori . . . . . . . . . . . . . . . . . . . . . . . . . . 16.59 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± δV and δI are the difference between these V , I determinations and those published by Hillenbrand (1997).. Audard et al.: The 2005 Outburst of V1118 Ori , Online Material p 3
Table A.2.
SMARTS average nightly magnitudes.MJD
U B V R I J H K · · · · · · · · · · · · · · · . ± .
07 11 . ± .
03 11 . ± . · · · · · · · · · · · · · · · . ± .
06 11 . ± .
03 11 . ± . . ± .
05 18 . ± .
03 17 . ± .
02 16 . ± .
01 14 . ± . · · · · · · · · · . ± . · · · · · · · · · · · · · · · · · · . ± . . ± .
08 18 . ± .
03 17 . ± .
02 16 . ± .
02 14 . ± . · · · · · · · · · . ± .
07 18 . ± .
03 17 . ± .
02 16 . ± .
01 14 . ± .
01 12 . ± .
07 11 . ± .
03 11 . ± . . ± .
03 17 . ± .
02 17 . ± .
01 15 . ± .
01 14 . ± . · · · · · · · · · · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 12 . ± .
05 11 . ± .
03 10 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
02 12 . ± .
01 12 . ± .
05 11 . ± .
03 10 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
05 11 . ± .
03 10 . ± . · · · . ± .
05 14 . ± .
02 13 . ± .
02 13 . ± . · · · · · · · · · · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
04 10 . ± .
02 10 . ± . · · · . ± .
02 13 . ± .
01 12 . ± .
01 11 . ± .
01 11 . ± .
03 10 . ± .
02 9 . ± . · · · . ± .
02 13 . ± .
01 12 . ± .
01 11 . ± .
01 11 . ± .
03 10 . ± .
02 10 . ± . · · · . ± .
02 13 . ± .
01 12 . ± .
01 11 . ± .
01 11 . ± .
03 10 . ± .
02 9 . ± . · · · . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
01 11 . ± .
03 10 . ± .
02 9 . ± . · · · . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
01 11 . ± .
03 10 . ± .
02 9 . ± . · · · . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
01 11 . ± .
03 10 . ± .
02 10 . ± . · · · . ± .
02 13 . ± .
01 12 . ± .
01 12 . ± .
01 11 . ± .
04 10 . ± .
02 9 . ± . · · · . ± .
05 13 . ± .
02 12 . ± .
02 11 . ± .
02 11 . ± .
05 10 . ± .
03 9 . ± . · · · . ± .
02 13 . ± .
02 12 . ± .
01 12 . ± .
01 11 . ± .
04 10 . ± .
02 10 . ± . · · · . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
01 11 . ± .
03 10 . ± .
02 9 . ± . · · · . ± .
02 13 . ± .
02 12 . ± .
01 11 . ± .
01 11 . ± .
04 10 . ± .
02 9 . ± . · · · . ± .
02 13 . ± .
02 12 . ± .
01 11 . ± .
01 11 . ± .
03 10 . ± .
02 10 . ± . · · · . ± .
02 13 . ± .
02 12 . ± .
01 11 . ± .
01 11 . ± .
03 10 . ± .
02 9 . ± . · · · . ± .
02 13 . ± .
02 12 . ± .
01 11 . ± .
01 11 . ± .
03 10 . ± .
02 9 . ± . · · · . ± .
03 13 . ± .
01 12 . ± .
01 11 . ± . · · · · · · · · · · · · . ± .
03 13 . ± .
01 12 . ± .
01 11 . ± . · · · · · · · · · · · · . ± .
03 13 . ± .
01 12 . ± .
01 11 . ± . · · · · · · · · · · · · . ± . · · · . ± .
01 11 . ± . · · · · · · · · · · · · . ± .
03 13 . ± .
01 12 . ± .
01 11 . ± . · · · · · · · · · · · · . ± .
04 13 . ± .
02 12 . ± .
01 11 . ± .
01 10 . ± .
03 10 . ± .
02 9 . ± . · · · . ± .
03 12 . ± .
03 12 . ± .
01 11 . ± .
01 10 . ± .
03 10 . ± .
02 9 . ± . · · · . ± .
11 13 . ± .
05 12 . ± .
04 11 . ± . · · · · · · · · · · · · . ± .
05 13 . ± .
02 12 . ± .
01 11 . ± . · · · · · · · · · · · · . ± .
05 13 . ± .
02 12 . ± .
01 11 . ± . · · · · · · · · · · · · . ± .
05 13 . ± .
02 12 . ± .
01 11 . ± . · · · · · · · · · · · · . ± .
04 13 . ± .
01 12 . ± .
01 11 . ± . · · · · · · · · · · · · . ± .
06 12 . ± .
03 12 . ± .
02 11 . ± .
01 11 . ± .
03 10 . ± .
01 9 . ± . · · · . ± .
03 13 . ± .
01 12 . ± .
01 11 . ± . · · · · · · · · · · · · . ± .
02 13 . ± .
02 12 . ± .
01 11 . ± .
01 11 . ± .
03 10 . ± .
01 9 . ± . · · · . ± .
02 12 . ± .
01 12 . ± .
01 11 . ± . · · · · · · · · · · · · . ± .
03 13 . ± .
01 12 . ± .
01 11 . ± . · · · · · · · · · · · · . ± .
02 12 . ± .
01 12 . ± .
01 11 . ± . · · · · · · · · · · · · . ± .
03 12 . ± .
02 12 . ± .
01 11 . ± .
01 10 . ± .
03 9 . ± .
01 9 . ± . · · · . ± .
02 13 . ± .
02 12 . ± .
01 11 . ± .
01 10 . ± .
03 10 . ± .
01 9 . ± . · · · . ± .
03 13 . ± .
03 12 . ± .
01 11 . ± .
01 10 . ± .
03 9 . ± .
01 9 . ± . · · · . ± .
02 13 . ± .
02 12 . ± .
01 11 . ± .
01 10 . ± .
03 9 . ± .
01 9 . ± . · · · . ± .
02 12 . ± .
02 12 . ± .
01 11 . ± .
01 10 . ± .
03 9 . ± .
01 9 . ± . · · · . ± .
02 12 . ± .
02 12 . ± .
01 11 . ± .
01 10 . ± .
03 9 . ± .
01 9 . ± . · · · . ± . · · · · · · · · · · · · . ± . · · · · · · · · · . ± .
02 12 . ± .
01 11 . ± .
01 10 . ± . · · · . ± . · · · . ± .
03 13 . ± .
02 12 . ± .
01 11 . ± .
01 10 . ± .
03 10 . ± .
02 9 . ± . · · · . ± .
03 13 . ± .
02 12 . ± .
01 11 . ± .
01 10 . ± .
03 9 . ± .
02 9 . ± . · · · . ± .
03 12 . ± .
02 12 . ± .
01 11 . ± .
01 10 . ± .
03 9 . ± .
01 9 . ± . · · · . ± .
14 13 . ± .
04 12 . ± .
02 11 . ± .
01 10 . ± .
03 10 . ± .
01 9 . ± . · · · · · · . ± . · · · · · · · · · · · · . ± . · · · . ± .
17 13 . ± .
01 12 . ± .
01 11 . ± .
01 10 . ± .
08 9 . ± . · · · · · · . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
01 10 . ± .
05 10 . ± .
02 9 . ± . · · · . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
01 11 . ± .
06 10 . ± .
02 9 . ± . · · · . ± .
01 13 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
04 10 . ± .
02 9 . ± . · · · . ± .
02 13 . ± .
01 12 . ± .
01 11 . ± .
01 11 . ± .
05 10 . ± .
02 9 . ± . · · · . ± .
02 13 . ± . · · · · · · · · · . ± .
03 9 . ± . · · · . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
01 11 . ± .
04 10 . ± .
02 9 . ± . · · · . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
01 10 . ± .
04 10 . ± .
02 9 . ± . · · · . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
01 11 . ± .
04 10 . ± .
02 10 . ± . . Audard et al.: The 2005 Outburst of V1118 Ori , Online Material p 4
Table A.2. continued.MJD
U B V R I J H K · · · . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
01 10 . ± .
04 10 . ± .
01 9 . ± . · · · . ± .
01 13 . ± .
01 12 . ± .
01 12 . ± .
01 11 . ± .
04 10 . ± .
02 9 . ± . · · · . ± .
01 13 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
05 10 . ± .
02 10 . ± . · · · . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
01 11 . ± .
04 10 . ± .
02 9 . ± . · · · . ± .
01 13 . ± .
01 12 . ± .
01 12 . ± .
01 11 . ± .
04 10 . ± .
02 9 . ± . · · · . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
01 11 . ± .
04 10 . ± .
02 9 . ± . · · · . ± .
01 13 . ± .
01 12 . ± .
01 12 . ± .
01 11 . ± .
04 10 . ± .
02 9 . ± . · · · . ± .
01 13 . ± .
01 12 . ± .
01 12 . ± .
01 11 . ± .
04 10 . ± .
02 9 . ± . · · · . ± .
01 13 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
05 10 . ± .
02 9 . ± . · · · . ± .
01 14 . ± . · · · . ± .
01 11 . ± .
07 10 . ± .
02 10 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± . · · · . ± . · · · · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
06 10 . ± .
02 10 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
05 10 . ± .
02 9 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
06 10 . ± .
02 10 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
06 10 . ± .
02 10 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
05 10 . ± .
02 9 . ± . · · · . ± .
02 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
05 10 . ± .
02 9 . ± . · · · . ± .
02 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
05 10 . ± .
02 10 . ± . · · · . ± .
02 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
05 10 . ± .
02 10 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
06 10 . ± .
02 10 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
05 10 . ± .
02 9 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
05 10 . ± .
02 9 . ± . · · · . ± .
01 13 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
05 10 . ± .
02 9 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
05 10 . ± .
02 9 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
05 10 . ± .
02 9 . ± . · · · . ± .
02 15 . ± .
02 14 . ± .
01 13 . ± .
01 12 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
02 15 . ± .
02 14 . ± .
01 13 . ± .
01 12 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
02 15 . ± .
02 14 . ± .
01 13 . ± .
01 11 . ± .
08 11 . ± .
03 10 . ± . · · · . ± .
02 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
06 11 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
06 11 . ± .
02 10 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
06 10 . ± .
02 9 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
06 10 . ± .
03 9 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
05 10 . ± .
02 9 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
05 10 . ± .
02 9 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
06 10 . ± .
02 10 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 13 . ± .
01 11 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 13 . ± .
01 12 . ± .
06 11 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
02 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
06 10 . ± .
02 10 . ± . · · · . ± .
02 15 . ± .
01 14 . ± .
01 13 . ± .
01 11 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
01 11 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
01 14 . ± .
01 13 . ± .
01 13 . ± .
01 11 . ± .
07 10 . ± .
03 10 . ± . · · · . ± .
01 14 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
06 10 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
07 11 . ± .
03 10 . ± . · · · · · · · · · · · · · · · . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
02 15 . ± .
02 14 . ± .
01 13 . ± .
01 12 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
02 15 . ± .
02 14 . ± .
01 13 . ± .
01 12 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
02 15 . ± .
02 14 . ± .
01 13 . ± .
01 12 . ± .
08 11 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
02 14 . ± .
01 13 . ± .
01 12 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
02 14 . ± .
01 13 . ± .
01 12 . ± .
08 11 . ± .
03 10 . ± . · · · . ± .
02 16 . ± . · · · · · · · · · . ± .
03 10 . ± . · · · · · · · · · . ± .
01 14 . ± .
02 12 . ± . · · · · · · · · · . ± .
02 15 . ± .
02 14 . ± .
01 13 . ± .
01 12 . ± .
08 11 . ± .
03 10 . ± . · · · . ± .
03 16 . ± .
02 14 . ± .
02 13 . ± .
02 12 . ± .
08 11 . ± .
03 10 . ± . · · · . ± .
02 16 . ± .
02 15 . ± .
01 14 . ± .
02 12 . ± .
08 11 . ± .
03 10 . ± . · · · . ± .
03 16 . ± .
02 14 . ± .
01 13 . ± .
02 12 . ± .
08 11 . ± .
03 10 . ± . · · · . ± .
02 16 . ± .
02 15 . ± .
01 13 . ± .
02 12 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
03 16 . ± .
03 15 . ± .
02 14 . ± .
02 12 . ± .
09 11 . ± .
03 10 . ± . · · · . ± .
02 16 . ± .
03 15 . ± .
02 14 . ± .
02 12 . ± .
08 11 . ± .
03 10 . ± . · · · . ± .
05 16 . ± .
05 15 . ± .
03 14 . ± .
02 12 . ± .
08 11 . ± .
03 10 . ± . · · · . ± .
03 16 . ± .
03 15 . ± .
02 14 . ± .
02 12 . ± .
08 11 . ± .
03 10 . ± . · · · . ± .
15 16 . ± .
09 15 . ± .
05 14 . ± .
04 12 . ± .
08 11 . ± .
03 10 . ± . . Audard et al.: The 2005 Outburst of V1118 Ori , Online Material p 5
Table A.2. continued.MJD
U B V R I J H K · · · . ± .
05 16 . ± .
07 15 . ± .
04 14 . ± .
04 12 . ± .
13 11 . ± .
04 10 . ± . · · · . ± .
29 17 . ± .
13 16 . ± .
08 14 . ± .
05 12 . ± .
10 11 . ± .
04 11 . ± . · · · . ± .
29 17 . ± .
28 16 . ± .
09 14 . ± .
05 12 . ± .
08 11 . ± .
03 10 . ± . · · · · · · · · · · · · . ± .
08 12 . ± .
18 11 . ± .
12 10 . ± . · · · . ± .
04 17 . ± .
03 16 . ± .
03 14 . ± .
02 12 . ± .
10 11 . ± .
04 11 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
08 11 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
09 11 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
08 11 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
09 11 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
07 11 . ± .
02 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
07 11 . ± .
02 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
08 11 . ± .
03 10 . ± . · · · . ± .
01 16 . ± .
01 15 . ± .
01 14 . ± .
01 11 . ± .
06 11 . ± .
03 10 . ± . · · · . ± .
01 16 . ± .
01 15 . ± .
01 14 . ± .
01 12 . ± .
07 11 . ± .
02 10 . ± . · · · . ± .
02 16 . ± .
01 15 . ± .
01 14 . ± .
01 12 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
02 16 . ± .
02 15 . ± .
02 14 . ± .
01 12 . ± .
09 11 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
07 11 . ± .
03 10 . ± . · · · . ± .
01 15 . ± .
01 14 . ± .
01 13 . ± .
01 12 . ± .
08 11 . ± .
03 10 . ± . · · · . ± .
04 16 . ± .
03 15 . ± .
03 14 . ± .
01 12 . ± .
09 11 . ± .
04 11 . ± . · · · . ± .
03 16 . ± .
02 15 . ± .
02 14 . ± .
01 12 . ± .
08 11 . ± .
03 11 . ± . · · · . ± .
04 17 . ± .
03 16 . ± .
03 14 . ± .
02 12 . ± .
09 11 . ± .
03 11 . ± . · · · . ± .
03 17 . ± .
02 16 . ± .
02 14 . ± .
01 12 . ± .
08 11 . ± .
03 11 . ± . . Audard et al.: The 2005 Outburst of V1118 Ori , Online Material p 6
Table A.3.
Villanova magnitudes.
MJD
V R I . ± .
05 12 . ± .
06 12 . ± . . ± .
02 12 . ± .
03 11 . ± . . ± .
07 12 . ± .
05 11 . ± . . ± .
06 12 . ± .
05 12 . ± . . ± .
06 12 . ± .
05 12 . ± . . ± .
03 12 . ± .
05 11 . ± . . ± .
33 12 . ± .
09 11 . ± . . ± .
07 12 . ± .
04 11 . ± . . ± .
02 11 . ± .
01 11 . ± . . ± .
02 11 . ± .
02 11 . ± . . ± .
03 12 . ± .
03 11 . ± . . ± .
03 12 . ± .
03 11 . ± . . ± .
04 12 . ± .
03 11 . ± . . ± .
03 13 . ± .
05 12 . ± . . ± .
06 12 . ± .
03 11 . ± . . ± .
04 12 . ± .
02 11 . ± . . ± .
04 12 . ± .
03 11 . ± . . ± .
02 12 . ± .
02 11 . ± . . ± .
03 13 . ± .
04 12 . ± . . ± .
09 13 . ± .
08 12 . ± . . ± .
13 13 . ± .
06 12 . ± . . ± .
12 13 . ± .
11 12 . ± . . ± .
22 14 . ± .
16 13 . ± . . ± .
34 14 . ± .
19 13 . ± . . ± .
16 13 . ± .
04 12 . ± . . ± .
04 13 . ± .
06 12 . ± . . ± .
08 13 . ± .
08 12 . ± . . ± .
08 14 . ± .
07 13 . ± . . ± .
11 13 . ± .
07 13 . ± . . ± .
12 13 . ± .
08 13 . ± . . ± .
07 14 . ± .
06 13 . ± . . ± .
14 14 . ± .
10 13 . ± . . ± .
08 14 . ± .
05 13 . ± . . ± .
10 14 . ± .
08 13 . ± . . ± .
08 14 . ± .
05 13 . ± . . Audard et al.: The 2005 Outburst of V1118 Ori , Online Material p 7
Appendix B: Details on the
Spitzer datareduction
B.1. IRAC
We used the MOsaicker and Point source EXtractor (MOPEX)software release of June 2007. We started from the BasicCalibrated Data (BCD) individual images to produce mosaicswith a (native) pixel size of 1 . ′′
22. For the data taken in sub-arraymode (PIDs 3716 and 41019), we collapsed the 3-dimensionalimage BCDs into 2-dimensional BCDs by using, for each pixel,the median of the 64 planes. For the uncertainty BCD files, weused the standard deviation of the 64 planes of the BCD file in-stead of using the input uncertainties. This method allows theremoval of most particle hits and the ingestion of the collapsedBCDs in MOPEX. We then used aperture photometry, centeredon V1118 Ori, using an extraction radius of 10 pixels (12 . ′′ µ m due to thestrength of the background intensity at this wavelength for radiilarger than 12.5 pixels. Indeed, the background is dominated byemission due to the nearby Herbig Ae star V372 Ori; this effectis less prominent in the other IRAC bands, but the presence ofthe Herbig star limited us to outer radii of 15 pixels). No aperturecorrection 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 detectoras input files for the MOPEX pipeline. A native pixel size of . ′′ was used to create the mosaic, and we used an extractionradius of 13 ′′ centered on V1118 Ori to derive the aperture pho-tometry. We used an annulus of radii 15 ′′ and 20 ′′ for the back-ground, and finally used an aperture correction factor of 1.17(from Table 3.12 of the MIPS data handbook, v. 3.3.1). We alsoinvestigated the background with a different method, i.e., to cal-culate the background in a nearby region using a circle with ra-dius 13 ′′ . Indeed, the MIPS24 background near V1118 Ori ishighly inhomogeneous. Notice that, while the uncertainty in theflux density is of the order 0.5 mJy, we estimate the true uncer-tainty, mainly due to the inhomogeneity of the background andthe difficulty to find a representative background region, to becloser to 15-20 mJy. B.3. IRS
We extracted the IRS spectra using the
Spitzer
IRS CustomExtraction (SPICE) v2.1.2 software and the post-BCD co-added nods For clarity, we explain the second technique for SL2 only in themain text. For the SL1 data, simply permute SL2 and SL1 in the text. and use the source-free region in the SL1 image for the back-ground of the SL1 nods. This method removes any contamina-tion by V1118 Ori in the SL1 image, but the downside is thatthe region of the sky covered by the SL1 slit (during the SL2observations) are away from V1118 Ori. Thus, the measuredbackground may not represent the background near V1118 Ori.A third technique consisted in using a background region nearV1118 Ori obtained during the observed nod. This technique al-lows to get a better estimate for the background near V1118 Ori,but has the disadvantage that V1118 Ori may contaminate thebackground, especially at longer wavelengths, since the stan-dard extraction width increases with increasing wavelength, totake into account the increasing size of the point spread functionof a point source. For the February 2005 observation, technique µ m (coming mostlyfrom the diffuse background emission) to check that the back-ground emission was adequately removed. For the March 2005observation, we used technique µ m in the SL1 slit. For thepost-outburst observations, we used technique λ > mic).For the high-resolution modules, we used the standard ap-proach to average the spectra from both nods using the full slitaperture. For the post-outburst spectra, we used the accompa-nying background spectra to obtain the background-subtractedspectra. This procedure worked nicely for the SH module (i.e.,up to . µ m), as demonstrated by the good subtraction ofthe PAH emission at 11.3 µ m. Notice that the H − S (1) λ . µ m line was apparently completely subtracted, whilethe H − S (2) λ . µ m line is faintly detected, Wealso detect [Ne II ] λ . µ m ( P / – P / ) an [S III ] λ . µ m ( P – P ) in the background-subtracted spec-trum (we will come back to this issue below). Furthermore, thesilicate feature and continuum flux up to µ m are consistentwith the SL flux, suggesting that the background subtractionwas not too far off. For the LH module, we subtracted only 95%of the background flux for the post-outburst data. Indeed, at thewavelengths covered by the LH, the background dominates theemission and is highly position-dependent. Although our back-ground pointing was close to V1118 Ori, there is considerableinhomogeneity in the diffuse emission of the Orion nebula, asobserved in the 8.0 and 24 µ m images. Figure B.2 shows theMIPS 24 µ m images together with the “field-of-views” of theIRS slits during the first outburst observation and in late 2008.We used the [Si II ] λ . µ m ( P / – P / ) as the proxy forbackground subtraction for the LH module, and noticed that ascaling factor of 0.95 was better than 1.0. We detect, however,excess flux in the [S III ] λ . µ m ( P – P ) line and faintexcess in the H − S (0) λ . µ m line (in contrast tothe S(1) line in the SH spectrum). The SH and LH spectra rel-atively well at 19.4 µ m, suggesting again that the backgroundsubtraction was adequate.We provide, however, a word of caution about the accuracyof the continuum flux level (especially above 30 µ m) and in thereality of the detected flux excess in some lines, especially in theLH wavelength region. Indeed, we have scaled the backgroundLH spectrum using the [Si II ] line. However, we remind thatV1118 Ori is located at the center of the Orion nebula, wherestrong ionization and excitation of the interstellar medium takesplace. Strong variations of the diffuse background emission takeplace, and they may differ for the continuum emission and for theline emission (mostly [Si II ] , [S III ], [Ne II ], and H ). Therefore, . Audard et al.: The 2005 Outburst of V1118 Ori , Online Material p 8 using a specific emission line to check that the background sub-traction is adequate may result in an incorrect subtraction of thecontinuum emission of the nebula and even in the subtractionof another emission line! Nevertheless, the similarities of thecontinuum shape in the low-resolution and high-resolution mod-ules (at least below 35 µ m) for the post-outburst data, the pres-ence of line excess in [S III ] in the LL spectrum and in both SHand LH spectra are suggestions that the background subtractionwas overall accurate. The presence of [Ne II ] in the SH modulebut not in the SL module demands, however, deeper analysis toconfirm the detection of the line. We degraded the SH and LHspectra down to SL and LL resolution (Fig. B.1, available onlineonly). The continuum shapes are well matched. But the degradedSH data covering the [Ne II ] line shows a strong line which is notcompatible with the SL data at this wavelength. This stronglysuggests that the SH background [Ne II ] line flux was too faintcompared to the line flux near V1118 Ori and that the measuredSH line flux is of background origin (from the Orion nebula orfrom diffuse emission close to V1118 Ori). In the case of H at12.28 µ m, the degraded SH spectrum is consistent with the SLspectrum: the contrast between the line and the continuum is toofaint to confirm the presence of this line in the SL spectrum. Thesame comment applies to the other molecular hydrogen line at28.2 µ m. Thus, we cannot confirm that this molecule is detectedin V1118 Ori’s spectrum. On the other hand the [S III ] lines inthe degraded SH and LH spectra have similar peak flux densi-ties as in the LL spectrum, indicating that these lines originatefrom the immediate vicinity of V1118 Ori. In summary, we be-lieve in the detection of the [S
III ] lines, while we have doubtsfor [Ne II ] and the H lines, but we provide nevertheless the linefluxes from the high-resolution module spectra in Table 3. In anycase, in view of Spitzer spatial resolution, higher spatial resolu-tion observations would be required to confirm the origin of theexcess line emission.We have further investigated whether the post-outburst back-ground observations could be used for the outburst SH and LHdata. In principle, the zodiacal light dominates the continuumbackground emission below about µ m, while the interstel-lar medium emission dominates above. The zodi contributionvaries as a function of time. We have used the Spitzer PlanningObservations Tool (SPOT) to determine the contribution at thethree different epochs of the IRS observations and found thatthey were of similar level (about 20-25 MJy sr − from 15 to24 µ m). For the SH spectrum, we preferred to use the post-outburst spectrum without scaling factor, since the PAH featureand the H line S(2) were well canceled (the S(1) line is slightlyoversubtracted). For the LH spectrum, we used a scaling factorequivalent to . , i.e., about the value of the ratio of the zodi-acal light contribution at µ m at in February-March 2005 andNovember 2008. This ratio also cancels out relatively well the[Si II ] line while the H S(0) line still remains detected. Withthis procedure, the resulting spectrum shows an increase of theflux with longer wavelengths, perhaps coming from an envelope,but this increase is not detected in the post-outburst data, cast-ing some doubt on the accuracy of the continuum shape in theSH and LH outburst spectra above 14 µ m. Since we have nooutburst LL data to confirm this shape, we prefer to err on thesafe side and claim that the high-resolution continuum shape for λ > µ m during the outburst is unreliable. The strengths ofthe emission lines may also be affected by improper backgroundreduction; indeed the flux levels are generally larger than in thepost-outburst spectrum (except for [Ne II ], but in this case, theline is probably not originating in the direct vicinity of V1118 Ori). Line fluxes from the high-resolution module data in out-burst are also given in Table 3. µ m)0.000.050.100.15 F l u x den s i t y ( Jy )
20 25 30 35 40Wavelength ( µ m)0.00.10.20.30.4 F l u x den s i t y ( Jy ) Fig. B.1.
Spitzer
IRS high-resolution post-outburst spectra (thickcurve) degraded to the resolution of the SL and LL modules. Thelatter data are shown as well (thin curve). Notice the good agree-ment in the continuum and for most lines, except for [Ne II ]. . Audard et al.: The 2005 Outburst of V1118 Ori , Online Material p 9
Fig. B.2.
Spitzer
MIPS 24 µ m image taken in March 2004 withoverlays of the IRS modules during the observations on 2005February 18 (top) and on 2008 November 14 after the outburst(bottom). Note that we only show the first position of the twonods, for clarity. Both nods for the background SH/LH obser-vations are shown in the bottom figure. A ′ scale is shown onboth North-orientated images. The image scale is linear from 0to 250 MJy sr −1