The 2008 outburst in the Young Stellar System Z CMa. III - Multi-epoch high-angular resolution images and spectra of the components in near-infrared
M. Bonnefoy, G. Chauvin, C. Dougados, A. Kospal, M. Benisty, G. Duchene, J. Bouvier, P. J. V. Garcia, E. Whelan, S. Antoniucci, L. Podio
AAstronomy & Astrophysics manuscript no. ZCMa˙bonnefoy˙aa c (cid:13)
ESO 2018October 17, 2018
The 2008 outburst in the Young Stellar System Z CMa
III - Multi-epoch high-angular resolution images and spectra of the componentsin near-infrared
M. Bonnefoy , , G. Chauvin , , C. Dougados , , ´A. K´osp´al , M. Benisty , , G. Duchˆene , , , J. Bouvier , , P. J. V.Garcia , E. Whelan , S. Antoniucci , and L. Podio Univ. Grenoble Alpes, IPAG, F-38000 Grenoble, France CNRS, IPAG, F-38000 Grenoble, Francee-mail: [email protected] Konkoly Observatory, Research Centre for Astronomy and Earth Sciences, Hungarian Academy of Sciences, PO Box 67, 1525Budapest, Hungary Astronomy Department, UC Berkeley, 501 Campbell Hall, Berkeley CA 94720-3411, USA Universidade do Porto, Faculdade de Engenharia, Rua Dr. Roberto Frias, s / n, P-4200-465 Porto, Portugal Department of Experimental Physics, National University of Ireland Maynooth, Maynooth, Co. Kildare, Ireland INAF - Osservatorio Astrofisico di Arcetri - L.go E. Fermi 5, I-50125 Firenze, ItalyReceived April 12, 2016 ; accepted June 22, 2016
ABSTRACT
Context.
Z CMa is a complex pre-main sequence binary with a current separation of 110 mas, known to consist of an FU Orionisstar (SE component) and an embedded Herbig Be star (NW component). Although it represents a well-studied and characterizedsystem, the origin of photometric variabilities, the component properties, and the physical configuration of the system remain mostlyunknown.
Aims.
Immediately when the late-2008 outburst of Z CMa was announced to the community, we initiated a high angular resolutionimaging campaign aimed at characterizing the outburst state of both components of the system in the near-infrared.
Methods.
We used the VLT / NACO and the Keck / NIRC2 near-infrared adaptive optics instrument to monitor the astrometric positionand the near-infrared photometry of the Z CMa components during the outburst phase and one year after. The VLT / SINFONI andKeck / OSIRIS integral field spectroscrographs were in addition used to characterize for the first time the resolved spectral propertiesof the FU Orionis and the Herbig Be component during and after the outburst.
Results.
We confirm that the NW star dominates the system flux in the 1.1-3.8 µ m range and is responsible for the photometricoutburst. We extract the first medium-resolution (R ∼ µ m) spectra of the individual components. TheSE component has a spectrum typical of FU Orionis objects. The NW component spectrum is characteristic of embedded outburstingprotostars and EX Or objects. It displays numerous emission lines whose intensity correlates with the system activity. In particular,we find a correlation between the Br γ equivalent width and the system brightness. The bluing of the continuum of the NW componentalong with the absolute flux and color-variation of the system during the outburst suggests that the outburst was caused by a complexinterplay between a variation of the extinction in the line of sight of the NW component on one hand, and the emission of shockedregions close to the NW component on the other. We confirm the recently reported wiggling of the SE component jet from [ FeII ]line emission. We find a point-like structure associated with a peak emission at 2.098 µ m coincidental with the clump or arm seen inbroadband polarization di ff erential imaging as well as additional di ff use emission along a PA = ◦ . The origin of these two structuresis unclear and deserves further investigation. Key words.
Techniques: high angular resolution, Binaries: general, Stars: pre-main sequence, individual (Z CMa)
1. Introduction
Z Canis Majoris (Z CMa) is a complex pre-main sequence binarymember of the CMa OB1 association (age < Send o ff print requests to : Micka¨el Bonnefoy: [email protected] authors were able to reproduce the UV-optical spectrum by amodel of an optically thick accretion disk surrounding an 1-3 M (cid:12) star with an accretion rate of 10 − M (cid:12) . yr − . The modelfailed to explain the near-infrared (NIR) part of the Z CMa spec-trum, however, as well as a strong increase in brightness in 1987.During the 1987 outburst, the optical spectrum was characterizedby a featureless and bluer continuum, a rise of Balmer lines, andemission lines of Fe II, Cr II, and Ti II (Hessman et al. 1991),which is unconsistent with typical FU Orionis activity.The problem was partly solved by Koresko et al. (1991) whodiscovered that Z CMa consists of a 110 mas binary. The SWcomponent of the pair is the FU Orionis star that dominates theoptical flux. The NW component dominates the NIR ( > µ m)to sub-mm spectrum, and the total luminosity. Follow-up obser- a r X i v : . [ a s t r o - ph . S R ] A ug . Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst vations (Whitney et al. 1993; Garcia et al. 1999; van den Anckeret al. 2004, hereafter VDA04) permitted the conclusion that thiscomponent is an Herbig Be star (hereafter HBe). VDA04 ob-tained a NIR (1.9–4.1 µ m) spectrum while the system was re-turning to its quiescent state. This spectrum shows Br α , Br γ ,and Pf γ in emission and the CO band-head in absorption around2.3 µ m. The authors suggested that the emission lines are formedinto an accreting circumstellar disk surrounding the HBe, andalso in an extended envelope above and below the disk plane.The emission line analysis led VDA04 to infer a B0IIIe spectraltype. Luminosity and e ff ective temperature estimates enables theauthors to derive a mass of 16 −
38 M (cid:12) for the HBe using variousevolutionary model predictions for ages of 3 . yr and 4 . yrrespectively.At large scales, this system reveals a rich outflow activity.A collimated optical outflow (3.9 pc; P.A. = ◦ ) and a bow-shock-shaped feature at 60 (cid:48)(cid:48) have been reported by Poetzel et al.(1989). Multiple outflow component profiles traced by opticalforbidden lines are present close to the source. Whelan et al.(2010) confirmed the existence of a jet driven by the HBe star(paper 2), which is unambiguously the driving source of theZ CMa parsec-scale outflow. The HBe star jet is seen oscillat-ing around a given position angle (jet wiggling), probably dueto a closer companion to the HBe star. A twin jet driven bythe FUOR component was also clearly identified. Millan-Gabetet al. (2001) also reported a cavity-like structure in adaptive op-tics J -band imaging, extending to the SW of the system, whichthey interpreted as light scattered o ff the wall of a jet-blown cav-ity aligned with the Z CMa large-scale outflows.Alonso-Albi et al. (2009) modeled the spectral energy distri-bution of the system by contributions from free-free emission, adisk of size 180 + − au tilted by 30 + − degrees, and an infallingenvelope (spherical or toroidal) surrounding each star that pos-sibly extends from 2000 to 5000 au and might be carved by theoutflows. The HBe itself is embedded in a dust cocoon (Whitneyet al. 1993). A sketch of the system can be found in Canovaset al. (2012). The dust distribution at the 500 au scale around ZCMa was investigated using NIR imaging polarimetry (Canovaset al. 2015; Liu et al. 2016). It reveals an extended filamen-tary structure (up to 2”) observed by Millan-Gabet & Monnier(2002), whose origin is unclear, as well as a polarized clumpcloser to the stars (0.3-0.5”).At the Astronomical Unit scale, the FUOR is strongly re-solved by the Keck-I interferometer to a level that is di ffi cultto explain with thermal emission of the accretion disk alone(Millan-Gabet et al. 2006). Hence its close environment appearsto be more complex than expected. The HBe component hasalso been observed with Keck-I (Monnier et al. 2005) and withAMBER on the VLTI (Li Causi et al. 2008; Benisty et al. 2010).Based on the Keck observations, the authors have modelledthe NIR emission as coming from a uniform ring. Nonetheless,they have finally concluded that the ring model was doubtfulfor this star, considering the large uncertainty of its spectraltype. Morover, the amount of infrared emission coming fromthis source alone at a given epoch is also poorly known, whichadds to the di ffi culty of interpreting the interferometric measure-ments.From January 2008 to October 2009, the system began to ex-perience an optical outburst, the largest reported in the past 90years of available observations (Grankin & Artemenko 2009).This outburst triggered additional observations of the system.Hinkley et al. (2013) presented adaptive optics JHKL band pho-tometry of the individual components. They concluded that theembedded HBe component is solely responsible for the outburst. Benisty et al. (2010) obtained spectrally resolved interferomet-ric observations of the HBe to study the hot gas emitting acrossthe Br γ emission line (paper 1). They found that the line profile,the astrometric signal, and the visibilities are consistent with theemission of a bipolar wind that may be partly seen through adisk hole inside the dust sublimation radius at the au scale. Theirmulti-epoch observations led them to suggest that the outburstis related to a period of strong mass-loss and not to a changein the extinction along the line of sight. The spectrophotomet-ric, spectropolarimetric, and polarimetric imaging observationsof the system (Szeifert et al. 2010; Canovas et al. 2012) suggestin contrast that the outburst was caused by changes in extinctionof the dust cocoon surrounding the HBe star.We present the results of a complementary high angular res-olution imaging campaign to resolve and characterize the prop-erties of each component of Z CMa in the NIR during and afterits 2008 outburst phase and to study its close environnement, inan attempt to better understand the origins of the bursts (extinc-tion, or accretion and ejection scenarios). Our new observationswere made from January to March 2009, that is, about one yearafter the start of the two-year outburst. We obtained an additionalobservation when the system returned to quiescence (see Fig. 1).We extracted the first medium-resolution (R ∼ µ m of each components of the system. Wepresent additional JHKL band photometric and astrometric dataof the binary. We report in Sect. 2 our observations and detail theassociated data reduction and analysis in Sect. 3. We analyze theevolution of the NIR photometric and spectroscopic propertiesof the components in Sects. 4 and Sect. 5, respectively. We dis-cuss these results in the context of other variable objects, and ofpast outbursts of the system in Sect. 6. We summarize our resultsin Sect. 7.
2. Observations
The Z CMa 2008 outburst was monitored with the NACO high-contrast adaptive optics (AO) imager of the VLT-UT4. TheNAOS AO system (Rousset et al. 2003) is equipped with a tip-tiltmirror, a 185 piezo actuator deformable mirror and two wave-front sensors (visible and IR). Attached to NAOS, CONICA(Lenzen et al. 2003) is the NIR (1 − µ m domain) imagingcamera equipped with a 1024 × ff erent epochs betweenJanuary 2009 and March 2009 (outburst phase). During the dif-ferent observing campaigns, the atmospheric conditions weresu ffi ciently stable to close the AO loop and resolve both compo-nents (see Fig. 2, Left panel ). The fainter FU Orionis componentin K -band is clearly visible on both Keck and VLT images. Bothimages were normalized in flux to show the flux variation thatis discernable between January and December 2009 in K band.The relative position of the FUOR with respect to that of theHBe component could be monitored well. The typical observingsequence included a set of five jittered images obtained using theJ, H, and K s bands with the S13 camera CONICA (mean platescale of 13.25 mas / pixel) and using the L’ filter with L27 (meanplate scale of 27.10 mas / pixel), leading to a total exposure timeof ∼
2. Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst V i s ua l M agn i t ude N I R SPE C S I N F O N I O S I R I SP A n t o . + AMBER V i s ua l M agn i t ude Fig. 1.
Light curve of Z CMa inferred from AAVSO observations. The spectroscopic observations of the HBe and FUOR are reportedin the figure. The dashed zone corresponds to the 2008 outbust of the system.
Table 1.
Observing log. Sr-2.2 µ m corresponds to the Strehl ratio measured at 2.2 µ m. UT Date Name Instr. Filter (Grism) Camera Exp. Time Sr-2.2 µ m Airmass Comment2009-01-31 Z CMa NACO J , H , K s , L (cid:48) S13, L27 90s (32%) 1.19HD54335 NACO J , H , K s , L (cid:48) S13, L27 90s 22% 1.25 psf-ref2009-02-06 Z CMa SINFONI J H K J H K J , H , K s , L (cid:48) S13, L27 90s (42%) 1.06HD54335 NACO J , H , K s , L (cid:48) S13, L27 90s 41% 1.11 psf-ref2009-03-11 Z CMa NACO J , H , K s , L (cid:48) S13, L27 90s (41%) 1.02HD54335 NACO J , H , K s , L (cid:48) S13, L27 90s 34% 1.04 psf-ref2009-12-07 Z CMa NIRC2 J cont , H cont , K cont / pix 8.6s - 1.29HIP33998 NIRC2 J cont , H cont , K cont / pix 8.6s 45% 1.51 psf-ref Fig. 2.
Left: VLT / NACO K s -band image of Z CMa on Januray31, 2009. Right: Keck / NIRC2 K cont image on December 7, 2009. On December 17, 2009 (transient phase), imaging observationswere also obtained using the AO system on the 10 m Keck II telescope (van Dam et al. 2004). We obtained direct imagesfrom 1.0 to 2.5 µ m that clearly resolved the two components ofthe system with the facility AO-dedicated NIR camera NIRC2,a 1024 × . ± . − scale and whose absolute orientation on the sky is0 . ± . ◦ (Ghez et al. 2008). We used the narrowband filters J cont , H cont , and K cont centered on 1.2132, 1.5804, and 2.2706 µ mto avoid the saturation of the star. The corresponding setups arealso reported in Table 1. Both components were easily resolvedto determine the relative flux and position (Fig. 2, Right panel ). We used the SINFONI instrument (Spectrograph for INfraredField Observations, see Eisenhauer et al. 2003a; Bonnet et al.2004), located at the Cassegrain focus of the VLT UT4 Yepun toobserve Z CMa during the outburst phase on February 6, 2009.The instrument provides AO-assisted integral field spectroscopy.It uses a modified version of the Multi-Applications Curvature
3. Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst
Fig. 3.
Illustration of the spectral deblending process at 1.65 µ m.Upper left: The initial Z CMa datacube where the two com-ponents are resolved. Upper right. During the observation, theadaptive-optics corrected PSF showed a strong astigmatism. Toachieve a proper extraction of the spectra of the Z CMa compo-nents, we duplicated the HBe star profile to create a PSF model.Lower left: The position and the flux of the individual sourceswere then retrieved using a modified version of the CLEAN al-gorithm. Lower right: The extraction error is estimated from theresiduals.Adaptive Optics system (MACAO, Bonnet et al. 2003) de-signed to feed the SPectrograph for Infrared Faint Field Imaging(SPIFFI, Eisenhauer et al. 2003b). SPIFFI pre-slit optics wherechosen to provide a spatial pixel scale of 12.5 ×
25 mas per pixel.Three di ff erent gratings were used to cover the J (1.1–1.4 µ m), H(1.45–1.85 µ m), and K band (1.95–2.45 µ m) at medium resolv-ing powers (2000, 3000, and 4000 respectively).The instrument was rotated to orient the binary horizontallyin the field of view (FoV) (see Fig. 2). Sky exposures wererecorded following a ABBA pattern. Additional o ff sets on theobject were chosen to increase the FOV to 850 ×
860 mas in theJ band, 850 ×
900 mas in the H band, and 950 ×
900 mas inthe K band. This also permitted us to artificially double the ver-tical spatial sampling. The source was bright enough at opticalwavelengths to be used as guide probe for the wavefront sensing.Hipparcos standard stars were also acquired soon after Z CMato correct the spectra for telluric features (see Table 1).
3. Data analysis
After cosmetic reductions (bad pixels, dark subtraction, flatfielding) using eclipse (Devillard 1997), we applied the deconvo-lution algorithm of Veran & Rigaut (1998) to obtain the positionof the HBe component relative to the FUOR at each epoch forboth NaCo and Keck data. The star HD54335 ( V = . K = .
6) and HIP33998 ( V = . K = . SINFONI data were reduced with the ESO data reductionpipeline version 1.9.8 (Abuter et al. 2006). We used in additioncustom routines to correct raw images from several electronic ef-fects that a ff ect the detector, as previously reported in Bonnefoyet al. (2014). The pipeline carried out cube reconstruction fromcorrected detector images. Hot and nonlinear pixels were taggedusing dark and flat-field exposures. Arc lamp frames acquiredduring the days following the observations enabled us to cal-ibrate the distortion, the wavelength scale, and the slitlet po-sitions. Slitlet distances were measured with north-south scan-ning of the detector illuminated with an optical fiber. Object-sky frame pairs were subtracted, flat-fielded and corrected forbad pixels and distortions. In each band, four datacubes were fi-nally reconstructed from clean science images and merged into amaster cube. The quality of the reduction was checked using theESO trending parameters. We obtained standard star datacubesin a similar way. However, their low signal to noise ratio (S / N)would have degraded the final Z CMa spectra, therefore we de-cided not to use it. We instead reduced datacubes of the telluricstandards HIP072241 (B3V) in the J band, HIP054970 (B5III)in the H band, and of HIP HIP071218 (G2V) in the K band ob-served at the airmass of Z CMa on di ff erent nights and with theinstrument pre-optics o ff ering a 0.25 arcsecond / spaxel sampling.The standard star spectra and the Z CMA spectra were smoothedto the same resolution. This change of resolution was estimatedfrom the arc lamp calibration files taken the same days as theobservations. Those telluric standards were found to provide thebest possible removal of the telluric lines.Although Z CMa components are resolved in our finaldatacubes, the two sources contaminate each other. We thereforedeveloped an improved version of the spectral extraction algo-rithm described in Bonnefoy et al. (2009) to retrieve deblendedspectra of the HBe and FUOR stars in each band (see Fig. 2).Our algorithm, CLEAN 3D has now been applied successfullyon di ff erent datasets (Bonnefoy et al. 2014; Bergfors et al.2016). We provide a reference concise description in thispaper. The positions of the two sources were initially fit witha Mo ff at function in each monochromatic cube slice. Theirvariation with wavelength caused by the atmospheric refractionwas fit with a low-order polynomial. A modified version ofthe CLEAN algorithm (Dumas et al. 2001) was then appliedon each slice to create a model of Z CMa while keeping theposition of the binary fixed. The standard star datacubes werenot reprentative of the PSF of Z CMa, which was a ff ected bya strong astigmatism and could therefore not be used as inputof the algorithm. We therefore chose to create a PSF templateduplicating the Herbig profile following the observed directionof the PSF lengthening. The position of the sources were thenrefined, an improved PSF model was created and a secondCLEAN step was performed on the original cube. The algorithmproduced two final datacubes containing the model of the HBeand the FUOR components. The typical error introduced by thespectral extraction and computed from residual maps rangesfrom 0.5 to 5%. Z CMa spectra were finally integrated inside
4. Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst K ( m ag ) H ( m ag ) J ( m ag ) JMD − K ( m ag ) H ( m ag ) J ( m ag ) JMD − Fig. 4.
Resolved Naco and Keck JHK photometry of FUOR component relative to the HBe compared to previous measurements(unresolved components) from the literature (VDA4, Koresko et al. 1991; Haas et al. 1993, and ref. therein). The epochs withsimultaneous JHK-band photometry are: [1] October 1990 (quiescence); [2-3-4], January 31, February 26, and March 11, 2009respectively (outburst); [5] December 2009 (returning to quiescence).
Table 2.
NaCo and Keck relative astrometry and photometry of FUOR relative to the HBe star in
JHK s UT Date ∆ PA platescale true North ∆ J ∆ H ∆ K s ∆ L (cid:48) (mas) (deg) (mas) (deg) (mag) (mag) (mag) (mag)09-01-31 111 ± . ± . . ± .
13 1 . ± .
10 2 . ± .
11 2 . ± . ± . ± . . ± .
04 0 . ± .
15 1 . ± .
04 1 . ± .
04 1 . ± . ± . ± . . ±
05 1 . ± .
05 2 . ± . ± . ± . ± ± − . ± .
20 0 . ± .
05 1 . ± . Table 3.
NaCo and Keck photometry of HBe, FUOR, and of the system in
JHK s UT Date J
HBe H HBe Ks HBe J FUOR H FUOR Ks FUOR J syst H syst Ks syst ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± circular appertures of 390 mas in the J band and 1250 mas in theH and K band. Spectra were divided by the standard star spectracorrected for intrinsic features and from their black body. Thestandard star spectra are found to provide a good correction oftelluric features in the J band and a moderate correction in the Hand K band. Spectra were finally scaled in flux by interpolatingthe NACO photometry of January 31, 2009 and February 26,2009 on the date of the SINFONI observations. We estimatethat the S / N measured on line-free regions on the Herbig spec-tra is 205 in the J band, 122 in the H band, and 120 in the K band.We repeated this procedure on the
Jbb (1.18–1.416 µ m)and Hbb (1.473–1.803 µ m) Keck / OSIRIS datacubes obtained onDecember 22, 2009 (Whelan et al. 2010) when the system hadreturned to its quiescent state. We also reduced and analyzed anadditional datacube obtained on the same night with the
Kbb (1.965–2.381 µ m) setting. The cubes have 20 ×
20 mas spax-els that sample the system separation well. The limited FoV of the cubes along the binary PA and residual artifacts limited theaccuracy of the extraction to ∼
10 %. We obtained the spectraof each component by integrating the flux of the sources overcircular apertures (R = Jbb and
Kbb datacubes, HD 109615 for the
Kbb dat-acubes). The standard star spectra were previously multiplied bya high-resolution (R = ∼ . The standard star was located at the edge ofthe FoV of the H-band cube, which introduced strong di ff eren-tial flux losses that translate into an additional slope in the finalspectra. We estimated and corrected the spectra for this slopeconsidering that the FUOR spectral slope should be identical tothe one of the SINFONI spectrum. The OSIRIS spectra have a http: // kurucz.harvard.edu / stars / VEGA /
5. Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst lower S / N than those of the SINFONI spectra (S / N =
31 in theJ band, 82 in the H band, and 70 in the K band). Finally, wenote that the K-band spectra are also slightly a ff ected by an os-cillating pattern (with a 0.08 µ m period) that might be related tonumerical artifacts found on the spaxels close to the edges of theFoV.
4. AO imaging of Z CMa
The photometry of the FUOR component relative the HBe ateach epoch for both NaCo and Keck data is reported in Table 3and compared in Fig. 4 with previous total and individual pho-tometry. From the literature, the FUOR is brighter in the quies-cent phase in J and H bands, whereas the HBe dominates in Kband. The system brightness also seems to decay significantlysince 1986, and presumably since the last FUOR burst. Duringthe 2008 outburst, our NaCo data show that the HBe componentbecame much brighter in all NIR wavelengths than the FUOR,as shown by a flux ratio inversion in J and H bands (see Table 2).This confirms that the HBe component is at the origin of the2008 outburst, in agreement with Hinkley et al. (2013). OurKeck observations, taken one year later, indicate that the systemdid not completely return to its quiescent phase in December2009 as the HBe component remains brighter in H-band.
In our 2009 NaCo observation in J or H bands, we did not re-detect any extended structures because the Keck data have com-paratively a higher S / N, as we did not intend to saturate our ob-servations to derive the photometry of the two components. Anextended structure at a projection angle (PA) compatible withthat reported by Millan-Gabet & Monnier (2002), Canovas et al.(2015), and Liu et al. (2016) is clearly seen in the L (cid:48) band (an-noted “A” in Fig. 5) after applying a radial profile subtraction.Its PA ranges from 160 to 250 ◦ and extends up to 2160 au.Two other structures (“B” and “C” in Fig. 5) are marginally de-tected in this band at PA ∼ ◦ . They are not detected atshorter NIR wavelengths or in polarimetry (Canovas et al. 2015;Liu et al. 2016). Their orientation does not coincide with the AC B
Diffraction spikes NE Detector pattern P A j e t HB e P A j e t F U O R Fig. 5.
VLT / NACO PSF-subtracted L (cid:48) -band image of Z CMa ob-tained on January 31, 2009. See Sect. 4.2 for details. PA ( deg r ee ) This paper0 2000 4000 6000 8000 10000 12000 14000JMD - 2445000 (day)90100110120 S epa r a t i on ( m a s ) Fig. 6.
Position in separation and position angle of the FUORcomponent relative to the HBe, based on the literature and ourNaCo and Keck observations.shocked emission regions seen in the wide-field images of ZCMa (Poetzel et al. 1989, Bouy et al. 2011, unpublished). Theycould be part of the cavity whose NW extension displays a com-plex structure (Poetzel et al. 1989) or be residuals from the stellarhalo subtraction.
The astrometry of the FUOR relative to the HBe was mea-sured at each epoch for our NaCo and Keck data. The resultsare reported in Table 2 and shown in Fig. 6. They complementolder epochs obtained using AO and speckle (Millan-Gabet et al.2001; Koresko et al. 1991; Haas et al. 1993; Barth et al. 1994;Thiebaut et al. 1995) and more recent measurements obtainedwith the polarimetric imaging mode of NaCo and SPHERE(Canovas et al. 2015, Antoniucci et al. 2016, in prep.). While theseparation does not seem to vary significantly, our astrometricmeasurements confirm a significant variation of position anglesince 1986. This might be consistent with a coplanar configura-tion between the binary and the HBe and FUOR disks (Malbetet al. 1993; Benisty et al. 2010) as often expected in binary sys-tems, although the jet orientations di ff er by at least 10 degreesbetween the two components (Whelan et al. 2010).
5. Spectro-imagery of the components of Z CMa
The flux-calibrated SINFONI spectra are shown in the 1.1-2.4 µ m range in Fig 7. The OSIRIS and SINFONI spectra of theFUOR and of the HBe are compared in each individual band inFigs. 8, 9, and 10. The 1.1-2.5 µ m spectrum of the HBe component is dominatedby a complex set of emission lines overlaid on a continuum witha red slope. These lines are decreased or absent in the OSIRISspectrum obtained during the quiescent phase. The NIR color-variation taken close in time with the OSIRIS spectra indicatesthat the continuum became redder when the system returned toquiescence. This is consistent with the slope variation of the con-tinuum in the J-, H-, and K-band spectra, although there are someuncertainties on the calibration of this slope in the OSIRIS data(see Sect. 5.2). The position of the emission lines emerging at
6. Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst
Fig. 7.
Flux-calibrated SINFONI spectra (from February 6, 2009) of the system (top), the HBe component (middle), and of theFUOR (bottom) in the 1.1-2.4 µ m range while the binary was in outburst. The lower resolution (R ∼
45) P1640 spectrum of thesystem components also acquired during the outburst phase (March 17, 2009, Hinkley et al. 2013) is shown in orange.more than 5% of the continuum flux are reported in Table 4.They were retrieved using an interactive least-squares multi-ple Gaussian fitting tool from the FUSE IDL library with anaccuracy of ∼ − ). The velocity of each iden-tified line, their corresponding transition, the associated non-dereddened fluxes, full-width at half maximum (FWHM), andequivalent widths (EW) for unblended lines complete Table 4.The velocities were corrected for the radial velocity of the sys-tem ( +
30 km.s − , Hartmann et al. 1989) and for the barycentric-to-helocientric velocity shift at the time of the observations (-11.2 km . s − ). The fluxes were estimated following the estima-tion and the removal of the continuum beneath the line usinga low-order Legendre polynomial. The associated error bar wasestimated from the RMS of the noise inside the continuum esti-mation zones surrounding the lines. FWHM are taken from theGaussian fit. Finally, equivalent widths (EW) and their associ-ated error bars are computed following the Sembach & Savage(1992) method. We propose an identification of most of the lines, for which werely on analogies between our spectrum and those of severalyoung embedded targets found in the literature (Kelly et al.1994; Greene & Lada 1996; Nisini et al. 2005; Gibb et al. 2006;Antoniucci et al. 2008; Connelley & Greene 2010; Covey et al.2011; K´osp´al et al. 2011b; Caratti o Garatti et al. 2013; Cooperet al. 2013). We completed and double-checked the identifica- The tool can be downloaded at http // fuse.pha.jhu.edu / analysis / fuse idl tools.html tion using three databases for atomic and ionic transitions , , .We also carefully flagged spurious lines that were introducedwhile dividing the spectrum by the telluric standards.We retrieved the same emission lines in the EX Or proto-type EX Lup captured during an outburst (K´osp´al et al. 2011b).The spectrum also resembles those of the outbursting proto-stars V1647 Ori, PTF 10nvg, and PV Cephei (Gibb et al. 2006;Covey et al. 2011; Hillenbrand et al. 2013; Caratti o Garatti et al.2013, see Fig. 11). We also find a good match with those ofclass I stars IRAS 03220 + + β represents the strongest line of the spectrum. In theH band, the more remarkable emission features are the series ofBrackett lines (order 10 to 22). The Brackett 23 line is expectedto be blended with a Mg I emission feature at 1.504 µ m. The Kband also has a moderate Br γ emission line (also present in theAMBER spectra of the HBe, Benisty et al. 2010).The J-band spectrum is also dominated by an O I line at1.1292 Å. The O I line is commonly attributed to fluorescenceexcitation by a UV continuum and / or to resonant absorption ofLy β photons (Kelly et al. 1994; Nisini et al. 2005). The non-detection of the O I line at 1.316 µ m favors the second hypothe-sis (K´osp´al et al. 2011b). http: // physics.nist.gov / PhysRefData / ASD / lines form.html http: // / ∼ peter / atomic / http: // / amp / ampdata / kurucz23 / sekur.html
7. Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst
Fig. 8.
J-band spectra of the Z CMa HBe ( top ) and FUOR( bottom ) components during the outburst phase (black and redline for the zoom; SINFONI spectrum from February 6, 2009)and in the quiescent stage (golden line; OSIRIS spectrum fromDecember 22, 2009). The fluxes of the OSIRIS spectra are nor-malized to the median flux of the SINFONI spectra between1.206 and 1.246 µ m and shifted to lower values for clarity.A forest of lines is also present from 1.15 to 1.22 µ m. Thisforest most likely corresponds to the broad feature seen in thespectra of HH100 IR and R Cra IRS2. It is composed of permit-ted lines of the neutral species of Fe, C, Si, Mg, and of inonizedCa. We also identify other lines of Fe I, Si I , C I, Al I, Na I, andMg I at longer wavelengths in the J, H, and K bands.Only one ionized species can be firmly identified in the spec-trum (Ca II). The first ionization potential of Ca (6.11 eV) islower than those of Si (8.15 eV) but remains comparable orhigher than those of Mg (7.64 eV), Al (5.98 eV) and Na I (5.14eV). The absence of Al II, Na II, and Mg II emission can be ex-plained by the energies of the upper transition levels (12.1-37.2eV) in the 1.1-2.5 µ m range. These energies are significantlyhigher than those involved for the Ca II (7.5 eV). As noted byK´osp´al et al. (2011b), the ionization potentials of the metallicspecies are lower (5.14-8.15 eV) than the one of carbon (11.26eV), oxygen (13.62 eV), and hydrogen (13.6 eV). This suggeststhat metals lies in a medium where hydrogen is mostly neutral.We report a weak [Fe II] emissions around 1.257 µ m, 1.534,and 1.644 µ m with several velocity components, associated witha micro jet studied in Whelan et al. (2010). The 1.644 µ m line ispartially blended with the Br line. We analyze structures asso-ciated with the 1.534 µ m line in Sect. 5.3.The spectrum has strong CO overtones of the X Σ + − X Σ + system seen in emission in the K band. These overtones havepreviously been reported by Hinkley et al. (2013), while the sys-tem was in the quiescence phase before the 2008 outburst. Thisis retrieved in many young stellar objects (e.g., Connelley &Greene 2010; Cooper et al. 2013). Connelley & Greene (2010)showed that Br γ coincides with the emergence of the CO band-head seen in emission. Since Br γ is a well-known accretiontracer (Muzerolle et al. 1998), the authors proposed that the COemission might arise from disk surfaces when the systems are Fig. 9.
H-band spectra of the Z CMa HBe ( top ) and FUOR( bottom ) components during the outburst phase (black line;SINFONI spectrum from February 6, 2009) and in the quies-cent stage (golden line; OSIRIS spectrum from December 22,2009). The fluxes of the OSIRIS spectra are normalized to themedian flux of the SINFONI spectra between 1.66 and 1.70 µ mand shifted to lower values for clarity.quite veiled and accretion rates are high. We follow this hy-pothesis and model them in Sect. 5.1.4. Alternatively, CO over-tones might be produced inside magnetospheric accretion fun-nels (Martin 1997). This line is not detected in the spectrum thatLiljestr¨om & Olofsson (1997) obtained while the system was inquiescence.To conclude, the comparison of the HBe spectrum to the oneof PTF 10nvg during its brighter stages suggests that we mayalso be seeing water-band emission in the spectrum of the HBeat 1.33-1.35 µ m, 1.49-1.55 µ m, and 1.70-1.75 µ m. Most of the lines have a FWHM greater than the instrumentalline-width (5.11 Å in the J band, 5.00 Å in the H band, 3.68Å inthe K band; measured on thorium-argon calibration lines). Theprofiles of the λ β , and of several other Brackettlines that are una ff ected by a strong blend (Br γ , Br 11, 17, 18,20, and 21) are reported in Fig. 12. They were subtracted fromtheir continuum and normalized to their highest value. The pro-files of the Pa β lines extracted from the SINFONI (black line)and OSIRIS (golden line) spectra are compared in the upper leftpanel. A Gaussian is fitted on the line core and is overlaid (grayline).The oxygen line core has a non-significant velocity. The linehas two broad asymmetrical wings that extend up to −
400 and + − . An additional unrelated feature (possibly an Al Iline) is present at − − .The resolved profiles of isolated H I emission lines (Pa β , Br11, Br 17, Br 18, Br 20, and Br 21) are shown in Fig. 12. ThePa β profile is divided into a main component characterized byasymmetrical wings and into a blueshifted lobe at − − .The lobe appears at lower velocities ( − − ) in the OSIRIS
8. Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst
Fig. 10.
Same as Fig. 9 but for the K band. The normal-ized NIRSPEC spectra (with respect to the continuum of theSINFONI spectra) of the system components obtained while thebinary was in a quiescent state (December 17, 2006) are overlaidin pink. The fluxes of the OSIRIS spectra are normalized to themedian flux of the SINFONI spectra between 2.18 and 2.22 µ mand shifted to lower values for clarity.spectrum and contributes more to the total line flux. We retrievedthese two components in the spectrum of EX Lup (K´osp´al et al.2011b), PV Cephei (Caratti o Garatti et al. 2013), and PTF 10nvg(Covey et al. 2011; Hillenbrand et al. 2013). The velocity isconsistent with an emission coming from the basis of the HBejet (Whelan et al. 2010). The line peak is slightly redshifted(42km.s − ), as is the one of PV Cephei (Caratti o Garatti et al.2013). The profile also has a red tail that extends up to ∼ − with a bump at + − . We retrieved this long tailin the H α and H β lines seen in the 1996 and 2000 spectra of ZCMa (van den Ancker et al. 2004). VDA04 proposed that the tailoriginates from the extended atmosphere of the HBe. The bumpis very consistent with the one seen in the spectrum of EX Lupduring outburst (K´osp´al et al. 2011b). In summary, the complexprofile of Paschen β (and related time variations) indicates thatthe line forms in (at least) two distinct regions: one at the basisof the ouflow, and another closer to the star.Brackett lines in the H band are slightly redshifted at thepeak ( + − ) and have a symmetrical profile extending to ± − . The Br γ profile is slightly asymmetric and red-shifted. This profile is coherent with the double-peaked asym-metric profile (peaks at ∼ −
40 and 120km.s − ) studied at R ∼ γ line originates from a bipolar wind atthe au scale. The same is expected for the other lines of the se-ries. We show in Fig. 13 the HI line ratios observed during the pho-tometric outburst phase for the HBe component. We only in-clude the HI Brackett lines detected in the H band to overcomethe e ff ect of poorly constrained reddening correction. The HI m m]1234 N o r m a li s ed f l u x + c on s t an t Z CMa (HBe; Feb. 6, 2009)PV Cephei (June 2012) V1647 ORI (March 2, 2005) PTF 10nvg (July 14, 2010) EX Lupi (July 24, 2008)1.55 1.60 1.65 1.70Wavelength [ m m]1.01.52.02.53.0 N o r m a li s ed f l u x + c on s t an t Z CMa (HBe; Feb. 6, 2009)PV Cephei (June 2012) V1647 ORI (March 2, 2005) PTF 10nvg (Sept. 23, 2010) EX Lupi (July 28, 2008)2.05 2.10 2.15 2.20 2.25 2.30Wavelength [ m m]1.01.52.02.5 N o r m a li s ed f l u x + c on s t an t Z CMa (HBe; Feb. 6, 2009)PV Cephei (June 2012) V1647 ORI (March 2, 2005) PTF 10nvg (July 14, 2010) EX Lupi (July 30, 2008)
Fig. 11.
Comparison of the NIR spectrum of the HBe compo-nent of Z CMa to spectra of variable young stellar objects PVCephei (Caratti o Garatti et al. 2013), V1647 ORI (Gibb et al.2006), PTF 10nvg (Covey et al. 2011; Hillenbrand et al. 2013),and EX Lupi (K´osp´al et al. 2011b). All spectra were smoothedto the lowest common resolution. Spectra of PV Cephei, V1647ORI, and PTF 10nvg were dereddened by A V =
8, 9, and6 mag, respectively to reproduce the Z CMa’s HBe pseudo-continuum slope. The H-band spectrum of EX Lupi is dered-dened by A V =
1, while the K-band spectrum is redenned by A V = -10 cm − , suggest-ing that the lines do not originate in the immediate vicinity of thestar. Such temperatures and densities could be found in the sur-face layers of the accretion disk, which are strongly irradiated bythe central star and accretion luminosity, or at the base of the out-flow. We derive a lower limit to the Pa β / Br γ ratio of 7.7 without
9. Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst
Fig. 12.
Profile of strong and weakly blended lines for the HBeand FUOR components. The black solid lines correspond to theSINFONI (outburst) data, while the golden lines correspond tothe OSIRIS spectra (post-outburst).
Fig. 13.
H band HI Brackett excitation diagram for the HBe ofZCMa (circles). All line fluxes are normalized to the HI Brackett11 line. Open circles show the lines that are probably contami-nated. Blue triangles show the line ratios detected in EX Lupi(K´osp´al et al. 2011b). The green curves are predictions fromcase B recombination from Hummer & Storey (1987) for T = K and n = cm − (lower curve) and for T =
500 K and n = cm − (upper curve). The red dotted lines show expected line ra-tios from blackbody emission at T = = K (upper curve).reddening correction. For the Pa β line, we integrated the emis-sion in the same velocity range as observed in Br γ , that is be-tween -300 and +
300 km s − . This very high ratio is marginallycompatible with case-B recombination and would exclude op-tically thick blackbody emission. The limit for optically thickblackbody emission is Pa β / Br γ ≤ β and Br γ lines originate from W avelength [ µ m ]- F (cid:104) [ - W . c m - . µ m - ] v = 2 (cid:65) (cid:65) (cid:65) (cid:65)
3T = 4350 K ; N CO = 2.0x10 c m - R o ut = 3.5 au Fig. 14.
CO bandhead seen in the SINFONI spectrum of the HBe(black) compared to our best -fit model for an outer radius of theline-emitting region of 3.5 au (red).the same regions as they show both distinct emission profiles(Sect. 5.1.2), and di ff erent behaviors with time (Sect. 5.1.5). Itis therefore likely that simple models assuming constant densityand temperature are not an adequate representation of the HI-emitting medium. Models of spherically expanding wind in LTEpredict both high Pa β / Br γ ratios and flat HI excitation diagrams(Lorenzetti et al. 2009; Antoniucci et al. 2011). The observedPa β flux and Pa β / Br γ ratio would require mass-loss rates ≥ − M (cid:12) yr − and an envelope thickness > R (cid:12) . Full radiative trans-fer modeling in either a disk or outflow model would be required,but this is beyond the scope of this paper. The CO overtones seen in emission in the HBe component arealso a characteristic feature of other EX Or (see Fig. 11). Eisneret al. (2014) have used interferometry to resolve the CO lineemitting areas around Herbig stars. They found that the regionsof the disks emitting in the CO bandhead lines are located be-tween 0.05 and 2 au from the central star. We therefore modeledthe CO lines of the HBe spectrum acquired in outburst follow-ing a similar approach as K´osp´al et al. (2011b) for EX Lupi. TheCO disk emission model and related fitting procedure are de-scribed in Appendix A. We conclude that the CO gas in Z CMahas an excitation temperature of about 4300 ±
100 K, and theCO column density is between 2 × cm − and 2 × cm − ,depending on the disk size. Combinations of the outer radius ofthe emitting disk region R out and of the CO column density N CO can give similar fit. We show in Fig. 14 the best-fitting model for R out = We report in Fig. 1 the dates of the spectroscopic observationsof the HBe on the light curve of the system.The low-resolution spectrum (R ∼
10. Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst the O I, the Paschen β emission line, and the Brackett linesseen in the SINFONI H-band spectrum of the HBe (Anto + ∼ β , Br γ ,and ν = −→ CO are reported in Table 6.The evolution of the Pa β line (Fig. 12) of the HBe sharessome behavior with the one of PV Cephei. The blueshifted lobethat is seen in absorption in the spectrum of PV Cephei (PCygni profile) only appears in the ourburst phase. The equiva-lent widths of the lines for the HBe and this object decreased by ∼
50% in 10 and 9 months, respectively, after the outburst. Thedecline of the equivalent width is consistent with the one of the ν = −→ CO overtone.In Fig. 15, we show that the equivalent width of the Br γ line decreases after the photometric outburst. This decline canbe reproduced by the function EW ( t ) = A / ( t − B ) β , where A and B are constant values modulating the amplitude and epochof maximum outburst, and β the decline factor. We used aLevenberg-Marquardt fitting tool to estimate A = − ± B = ±
25, and β = . ± .
20. The B parameter corre-sponds to an epoch of maximum outburst on August 30, 2008.Unfortunately, this falls at a time when the V band was not mon-itored, but the system was considered to be in outburst at thattime. We caution of course that the analysis is limited by thesmall number of epochs. The equivalent widths of the line de-cline with that of the V-band magnitude of the system (correla-tion factor = (Kafka, S., 2016). This line is not detected in the November 1992low-resolution spectrum of the HBe, which confirms that it is re-lated to the system activity.The NIRSPEC spectrum obtained while the system was ina quiescent state shows a He I line (2.1125 µ m, 3p P ◦ -4s S).This line is not retrieved in the SINFONI and OSIRIS spectra. Itis seen in some YSO spectra and is commonly associated withthe stronger 2.0508 µ m He I line (2s S-2p P ◦ ) (e.g., Cooper et al.2013; Murakawa et al. 2013) which has been detected by van denAncker et al. (2004) in the IRTF spectrum of the system takenwhen it was returning to quiescence (post-2000 outburst). Thissecond line is located outside of the NIRSPEC coverage, but it isstill missing from our SINFONI and OSIRIS spectra of the HBe.Cooper et al. (2013) proposed that these lines form by collision-induced excitation in a wind. The disppearance of this line sys-tem in the outburst (SINFONI) and post-outburst (OSIRIS) statemight be due to the evolving contrast with the continuum emis-sion, or to the variable extinction in the line of sight (see below).The NIR broadband photometry is mostly sensitive tothe continuum emission. Therefore, the decrease of the line-equivalent width coupled to the stagnation (K band) or decrease(V band, J and H bands) of the continuum flux indicates that theflux produced by the excited regions does decline as well or thatthe absorption increases. The bluing of the slope of the HBe inthe outburst phase supports the hypothesis of a decrease in ex-tinction as well as an increase of the internal heating of the inner http: // portions of the disk due to an increase in accretion rate. We notethat the continuum slope of the SINFONI spectrum is consis-tent with the one obtained with P1640. Both the SINFONI andOSIRIS K-band spectral continua look redder than the one ob-tained in November 1992 (Fig. 1 of Liljestr¨om & Olofsson 1997)while the system was in quiescence.The evolution of the HBe in color-color and color-magnitudediagrams (Fig. 16) indicates that the photometry of the star hasone component that follows the interstellar reddening vector.This suggests that the outburst is at least partly due to a reducedextinction in the line of sight. van den Ancker et al. (2004) re-ported a spectral type B0IIIe for the HBe. Assuming that thesestars have NIR photometric colors close to 0, the location ofthe HBe in the Fig. 16 indicates that the star was extincted byA v =
12 mag during the 2008 outburst. This is consistent withA v =
10 mag measured by Hinkley et al. (2013) from the SEDanalysis of the star obtained during the outburst. This also agreeswith the lower limit on A v derived from line ratio in Sect. 5.1.3.The reduced A v found by Hinkley et al. (2013) is also con-sistent with the higher flux level of their spectra compared toSINFONI’s and shown in Fig. 7, although we cannot excludethat the di ff erent flux levels may come from di ff erences in theflux-calibration methods used the two sets of spectra.The variation in H-K color of the HBe component showsan additional contribution perpendicular to the reddening vec-tor. This deviation is reproduced by several EX Ors (Lorenzettiet al. 2006; K´osp´al et al. 2011a; Lorenzetti et al. 2012), and inparticular EX Lup. EX Lup represents the NIR spectrum of theHBe well (Fig. 11). Lorenzetti et al. (2012) showed that duringphotometric outbursts of EX Ors, an additional blackbody com-ponent appears with temperatures between 1000 and 4500 K anda flux corresponding to blackbody emission from a uniform diskof radius between 0.01 and 0.1 au. Hinkley et al. (2013) fit theSED of the HBe with two blackbodies at T e ff = A V =
10. But their analysis relies on photometric pointsobtained while the system was at di ff erent stages. As shown inFig. 3 of Giannini et al. (2016b), an increase in temperature ofthe blackbody component in the SED of the HBe during the out-burst phase might well explain the variation in colors perpendic-ular to the interstellar reddening vector. The NIR spectrum of the FUOR component of Z CMa is char-acteristic of M8-M9 giants of the IRTF spectral library (Rayneret al. 2009), as expected for FU Orionis objects (see Hartmann& Kenyon 1996) with a luminosity in NIR totally dominatedby the accretion disk flux and with the cooler regions of thedisk producing deep molecular absorptions. Strong broad ab-sorption bands of H O, TiO, VO, and CO lines are thereforedetected. The CO lines in absorption of the FUOR were previ-ously marginally detected in the unresolved spectrum of Z CMaby Reipurth & Aspin (1997), through the strong contributionof the HBe component the K-band flux. These absorptions aretypical of FU Orionis objects (Mould et al. 1978, Carr et al.1989). The individual spectrum of the component obtained byLiljestr¨om & Olofsson (1997) and Hinkley et al. (2013) showedthat the CO in absorption can be associated with this component.The SINFONI (outburst) and OSIRIS (returning to quiescence)spectra shown in Figs. 8, 9, and 10 unambiguously confirm thisresult.The OSIRIS and SINFONI spectra display a strong Pa β ab-sorption. This line is retrieved in the spectra of FU Orionis ob-jects HBC687 ( = IRAS 19266 + = IRAS
11. Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst E qu i v a l en t w i d t h [ A ng s t r o m s ] OUTBURST QUIESCENCEAMBERAMBERSINFONISINFONIOSIRISOSIRIS4700 4800 4900 5000 5100 5200 5300MJD - 2 450 000 [Days]0-1-2-3-4-5 E qu i v a l en t w i d t h [ A ng s t r o m s ] l og (- x E qu i v a l en t w i d t h ) AMBERAMBERSINFONISINFONIOSIRISOSIRIS
Fig. 15.
Decline of the Br γ equivalent width with time (left) and with the system V-band magnitude (right). -1 0 1 2 3 4[H-K]-0.50.00.51.01.52.0 [ J - H ] Herbig123 4 5 Fu Ori1 23 4 5A V =0A V =1A V =5A V =10A V =15 1 - Oct. 6, 19902 - Jan. 31, 20093 - Feb. 26, 20094 - March 11, 20095 - Dec. 07, 2009 0 1 2 3 4J-H (mag)681012 J ( m ag ) A V =0A V =1A V =5 A V =10A V =15Herbig 1234 5Fu Ori 12 345 5 - Dec. 07, 20094 - March 11, 20093 - Feb. 26, 20092 - Jan. 31, 20091 - Oct. 6, 1990 0 1 2 3 4 5H-K (mag)46810 H ( m ag ) A V =0A V =1A V =5A V =10A V =15Herbig 123 4 5Fu Ori1 234 5 5 - Dec. 07, 20094 - March 11, 20093 - Feb. 26, 20092 - Jan. 31, 20091 - Oct. 6, 1990 Fig. 16.
Evolution of the FUOR and HBe NIR colors and photometry when Z CMa was in a quiescent state (1; October 1990observations), in outburst (2-3-4; encircled disks; corresponding to NaCo observations on January 31, February 26, and March 11,2009 respectively) and returning to a quiescent state in December 2009 (5). The reddening vector for visual extinctions of 0, 1,5,10,and 15 mag as well as the colors of pre-main-sequence stars (solid line, Pecaut & Mamajek 2013) and the locus of CTTS (dashedline, Meyer et al. 1997) are overlaid for comparison.21454 + β absorption can be fit by twoGaussian functions with FWHM of 8 and 11.5 Å (broadly con-sistent with the spectral resolution), intensity ratio beneath thepseudo-continuum of 1.5, and velocities of −
265 km.s − and −
102 km.s − , respectively. This double-peaked profile is char-acteristic of FU Orionis stars in the optical (e.g., Hartmann &Kenyon 1985). It is reminiscent of the H γ and H δ line profiles ex-tracted from the optical spectrum of Z CMa acquired in February1983, while the system was not in outburst (Covino et al. 1984).The spectra of HBC687 and V1735 Cyg seem to exhibit an asy-metric Pa β line profile similar to the one of Z CMa, although thewavelength sampling of these comparison spectra is lower thanSINFONI’s. The line profile of the FUOR does not change sig-nificantly during the one-year lag corresponding to the OSIRIS and SINFONI spectra. The line properties are clearly not com-patible with the hypothesis of an unresolved binary.The colors of the FUOR recorded in 1990 and during the2008 outburst (Fig. 16) are redder than those of pre-main-sequence stars earlier than K4 (J-H = = ffi cult to relate thisbehavior to those of other FU Ors since there has not been anextensive follow-up of the NIR photometry of these objects sofar. Our observations are consistent with the variation in opticalcolors, in the period which show a progressive bluing of Z CMacolors consistent with the behavior of the HBe in the NIR, but
12. Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst
Table 7.
Line identification in the 1 . − . µ m spectrum of theFUOR components λ obs Element Transition( µ m)1.10399 TiO 0–0 band of Φ ( b Π − d Σ )1.11752 TiO 1–1 band of Φ ( b Π − d Σ )1.13502 TR . . . . . . A Π − X Σ − Φ ( b Π − d Σ )1.25341 TR . . . a D / − a D / Φ ( b Π − d Σ )1.27870 TR . . . β ) 3 ∗ − ∗ . . . O 2 ν , ν + ν , 2 ν ,2 ν + ν , ν + ν . . . . . . . . . CO ? 5–2 band of X Σ + − X Σ + CO ? 6–3 band of X Σ + − X Σ + . . . O ν + ν , ν + ν , 3 ν . . . . . . . . . . . . . . . CO 2–0 band of X Σ + − X Σ + CO 3–1 band of X Σ + − X Σ + . . . O ν , ν , 2 ν also another dependence that could be induced by the variabilityof the FUOR component (Grankin & Artemenko 2009).The intrinsic variation of the FUOR colors could translateinto NIR slope variation of our continuum. In that case, the slopeof the spectral continuum of the HBe in the H-band OSIRIS datacould be uncertain because it has been based on the hypothesisof a non-variation of the spectrum of the FUOR (see Sect. 3.2).This does not change our conclusions on the variation of the HBecontinuum slope, which mostly relies on the broadband NIR col-ors reported in Table 3. The spatial sampling of the SINFONI cubes enable us to lookfor extended structures down to ∼ ∼ ◦ and 245 ◦ , respectively into 1.257 µ m and 1.644 µ m[FeII] lines. In a recent analysis of H α and [ OI ] (655.6 nm and629.5 nm) AO-imaging data (Antoniucci et al. 2016, submittedto A&A), we report the detection of the wiggling of the FUORjet. We chose here to re-investigate our data to look for structuresin other emission lines in the context of the results of Antoniucciet al. 2016.We report extended emission at 1.200, 1.205, and 1.533, and2.098 µ m following the subtraction of the continuum emission(see Appendix B for the details). We classify these emissionsinto three categories: – The 1.533 µ m emission corresponds to the [ FeII ] line seenin our spectrum of the system (Table 4 and Fig. 9). The emis-sion map at more than 3 σ the noise level is shown in Fig. 17.The dashed ellipse corresponds to the era within which thecontinuum substraction leaves strong residuals. The emis-sion is associated with the FUOR micro jet (orange dashedline in the figures). It is elongated along a position angle of ∼ ◦ . The emission is detected in the same velocity rangeat 1.257 and 1.644 µ m. The emission has a sinusoid shapewith a ∼
75 mas semi-period from −
400 to −
100 km.s − thatis consistent with the wiggling found by Antoniucci et al.2016. – The 1.200 and 1.205 µ m line emissions may be associatedwith [ FeII ] lines at 1.2000278 µ m and 1.2054490 µ m, respec-tively. We assume that we correctly identified the lines tomap the emission in the velocity space in Fig. 18. Theselines trace an elongated emission at a position angle of ∼ ◦ (green dashed lines in the Figs. 17 to 18). The emission isseen from -600 to 110 km.s − . Its elongation is in the samedirection as the K / SPHERE (
N CntJ filter) may help to clarify the natureof this structure. – A clump is detected at 2.098 µ m (Fig. 18) at a PA of 164 ± ◦ and separation of 313 ±
13 mas (291-360au). This clump isat the same position as the polarized clump or arm detectedin the H (1.66 µ m) and Ks (2.18 µ m) band by Canovas et al.(2015) and Liu et al. (2016). The He I and H emission linesseen in jets and embedded objects in this wavelength rangetranslate into an absolute velocity (cid:29) − . Plausibleexplanations are that the line is incorrectly identified, or thatthe structure is unrelated to the system. We detect anotherclump at a position angle of ∼ ◦ (labeled as ? in the Fig.18). This additional clump is also found at 2.11 µ m, and couldcorrespond to an emission in the He I line (2.112583 µ m)blueshifted by −
400 to −
200 km.s − . Nonetheless, it fallsclose to the PSF Airy ring of the FUOR which moves withwavelength and might produce some residual emission at thecontinuum subtraction step.We were unable to detect any extended structures apart fromthe probable di ff raction spikes associated with the HBe andFUOR point sources in images derived from the cubes collapsedin wavelengths and with the stellar halo of each component re-moved with a Gaussian smoothing (with a FWHM =
5, 10, and15 pixels).
6. Origins of the outburst: extinction or accretion?
The architecture of the Z CMa system is reminiscent of thoseof other EX Ors objects with binary companions such as V1118Ori (Reipurth et al. 2007), VY Tau (Leinert et al. 1993), andXZ Tau (Hartigan & Kenyon 2003), altough Z CMa may be themost massive system of this class of outbursting binaries, and This value is consistent with the interval given in Poetzel et al.(1989) for the large-scale jet and by Whelan et al. (2010) andAntoniucci et al. (2016) for the FUOR micro jet. 13. Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst O ff s e t [ m a s ] (a) -400 to -340 km.s -1 NE FUOR HBe 0 200 400 600 800Offset [mas]0200400600800 O ff s e t [ m a s ] (b) -300 to -250 km.s -1 NE FUOR HBe 0 200 400 600 800Offset [mas]0200400600800 O ff s e t [ m a s ] (c) -230 to -100 km.s -1 NE FUOR HBe 0 200 400 600 800Offset [mas]0200400600800 O ff s e t [ m a s ] (d) -100 to +110 km.s -1 NE FUOR HBe
Fig. 17.
Extended structures observed in the [
FeII ] line at 1.53 µ m O ff s e t [ m a s ] m m -600 to 110 km.s -1 NE FUOR HBe 0 200 400 600 800Offset [mas]0200400600800 O ff s e t [ m a s ] m m -600 to 110 km.s -1 NE FUOR HBe 0 200 400 600 800Offset [mas]0200400600800 O ff s e t [ m a s ] m m NEFUOR HBeClump? Fig. 18.
Extended structures observed at 1.200, 1.205, and 2.098 µ mthe only one with a FU Or companion. These three systems haveprojected separations of 72 au, 95 au, and 43 au respectively. It ispossible that the outbursts of Z CMa are related to instabilities inthe disk surrounding the HBe at the corotation radius (e.g., muchcloser to the star) where gas accumulates until it is accreted dur-ing outbursts (D’Angelo & Spruit 2012). This scenario may ex-plain the behavior of EX Lup (D’Angelo & Spruit 2012; Banzattiet al. 2015). In this case, are our spectrophotometric data of theHBe of Z CMa more compatible with the scenario of an accre-tion + ejection event, or a variation of the extinction in the line ofsight?The fact that the jet of the HBe wiggles (Whelan et al. 2010),the existence of knots along the direction of the HBe outflow,and the change in velocity of the blueside wing of the Paschen β line associated with the HBe outflow suggest that the 2008outburst and spectrophotometric properties of the most massivestar of the system are related to an accretion + ejection event. Thespectrum of the HBe, quite similar to the one of the prototypeEX Lup, strengthens this conclusion. Nonetheless, because (i)the Br γ line emission may originate at least partly in a disk wind(Benisty et al. 2010) and (ii) the empirical relation between theBr γ flux and the accretion rate seems to break down in the Beregime (Donehew & Brittain 2011), the accretion rate as wellas the extinctions from line ratio (Pa β / Br γ ) cannot be derivedreliably from our NIR spectra.Szeifert et al. (2010) claimed that the 2008 outburst is relatedto a change in extinction from the dust cocoon surrounding theHBe based on the level of optical polarization in both continuumand spectral lines along a position angle roughly perpendicularto the outflow launched by the HBe. The characteristics of theHBe component during the 2008 outburst are not fully typicalof EX Or objects and outbursting stars. The variation in the NIRcontinuum slope (colors) of the HBe with respect to the absoluteNIR flux is anticorrelated with the one of V1647 Ori, PV Cep, or V1118 Ori (Gibb et al. 2006; Lorenzetti et al. 2009; Gianniniet al. 2016a). The emission lines of these four reference objectsare correlated with their overall NIR brightness (outburst or qui-escence) like Z CMa, based on the current stage of observations(temporal baseline, frequency of the measurements) of all thosesystems. Stelzer et al. (2009) noted that V1118 Ori had a di ff er-ent X-ray luminosity and temperature variation during an out-burst with respect to the behavior of Z CMa during the 2008 out-burst, thus pointing to a di ff erent outburst mechanism for thesesystems. Z CMa may resemble the case of PTF 10nvg where thespectrophotometry of the system seems to be driven by the ro-tating circumstellar disk material located at short separation ( < + ejection event modulatedby absorption in the line of sight during the quiescent stage of2010 (OSIRIS spectra). This would explain why the CO band-head is seen in emission during the 2006 quiescent stage (the lastone before the 2008 outburst) and then could not be seen at allin the OSIRIS spectrum taken during the 2009 quiescent period(veiling).
14. Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst
7. Conclusion
We have obtained multi-epoch NIR astrometric, photometric,and spectroscopic data of the young system Z CMa usingadaptive-optics-fed instruments. Our data enabled us to ob-tain resolved 1-4 µ m photometry as well as medium-resolution(R ∼ + ejection changesand variable extinction in the line of sight. We find in particulara correlation between the strength of the Br γ emission and thebrightness of the system. Nevertheless, we cannot have a directaccess to the reddenning and accretion rate because the emissionlines that would enable for an estimation originate in di ff erentlocations around the Herbig objet (disk, outflow). This preventsus from firmly linking the spectrophotometric properties of thesystem to their origins. The FU Or companion also experiencedcolor variations at a constant luminosity during the outburst thatdo not seem to be related to variable extinction.We detected the extended emission the FUOR jet at 1.53 µ m ([Fe II] line). We also resolved an extended emission at1.200 and 1.205 µ m that extends along a position angle of 214 ◦ ,and whose origin is unclear. To conclude, we identified a pointsource at 2.098 µ m that is concomitant with a polarized clumpseen in multiple NIR imaging data.The system is undergoing a new outburst in 2016. Similarobservations as conducted in our study, but with a better tem-poral sampling, coupled to observations at longer wavelengthswith the mid-infrared instrument VISIR (Lagage et al. 2004)and ALMA may dramatically improve our understanding ofthe architecture of the system, and how it relates to the accre-tion + ejection diagnostics used in the NIR. Appendix A: Details on modeling the CO emission
We first considered a disk defined by inner and outer radii ( R in , R out ) and an inclination ( i ). The radial velocity was computedfor each point of the disk assuming that the disk material is inKeplerian rotation around a star with a mass M (cid:63) . As for the COemission, we used Eq. 10 from Kraus et al. (2000) to calculatethe absorption coe ffi cient κ . A simple Gaussian was assumed forthe spectral function. This accounts for the thermal and turbulentbroadening and for a medium-resolution spectrum. κ was multi-pled with the column density of CO to obtain the optical depth τ . We assumed a simple slab, so that the column density and ex-citation temperature of CO ( T ex ) were identical at every locationin the disk. We calculated the intensity I following: I = BB × (1 − e ( − τ/ cos( i )) ) (A.1)where BB is the Planck function corresponding to the excitationtemperature.The computed spectrum was redshifted (or blueshifted) bythe radial velocity appropriate for each point of the disk. Wefinally integrated over the whole disk in each spectral channeland smoothed the resulting spectrum to account for the finiteinstrumental spectral resolution of SINFONI.We varied T ext between 1500 K and 4500 K, in steps of150 K, and a column density N CO between 10 . cm − and10 . cm − in logarithmic scale, in steps of 0.1 in the exponent. We took R in = i = ◦ , and M (cid:63) = (cid:12) . We considered R out from 0.5 to 3.5 au, with 0.5 au increments.We compared the first two overtones of the HBe (followingthe removal of the continuum in the SINFONI spectrum) to thegrid of 3381 model spectra using a χ minimization. Appendix B: Removal of the continuum emission
The removal of the spectral continuum in the SINFONI dat-acubes is required to look for faint structures close to the bi-nary. This is usually done by fitting spaxel to spaxel the spectralcontinuum around the object line. Here, the intrinsically di ff er-ent continua from the FUOR and HBe made this removal morecomplicated. We found that using neighboring cube slices as ref-erence PSF provided a good solution to remove most of the con-tinuum flux of each component. The technique works better thanthe subtraction in the spectral range (fitting a low-order polyno-mial) or the spatial subtraction of a PSF model of each compo-nent using the extracted cubes derived from the CLEAN 3D algo-rithm (see Sect. 3.2).
Acknowledgements.
We would like to thank particularly the sta ff of ESO-VLTfor their support at the telescope as well as Sasha Hinkley and Mario VanDen Acker, who kindly provided their NIR photometry and spectra of the ZCMa system. We are very grateful to Mike Connelley, Kevin Covey, LynneHillenbrand, Erika Gibb, and Alessio Caratti o Garatti for providing their spec-tra of outbursting and embedded objects. We acknowledge with thanks the vari-able star observations from the AAVSO International Database contributed byobservers worldwide (Eric Blown in particular) and used in this research. Thiswork was supported by the Momentum grant of the MTA CSFK Lend¨ulet DiskResearch Group. We acknowledge partial financial support from the
ProgrammesNationaux de Plan´etologie et de Physique Stellaire (PNP & PNPS) and the
Agence Nationale de la Recherche , in France.
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16. Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst
Table 4.
Line identification in the 1 . − . µ m spectrum of the HBe components λ obs Element Transition ∆ V FWHM Flux EW Remarks( µ m) (km . s − ) (Å) (10 − erg . s − cm − ) (Å)1.11306 Fe II 4 p F ◦ / − s G /
42 13.6 1 . ± . − . ± .
12 CP, Blend with Fe I1.12920 O I 3 p P − d D ◦
42 10.8 4 . ± . − . ± .
16 CP1.13370 C I ? 3 p P − d D ◦ . . . . . . . . . SB, or TR?1.13502 TR . . . . . . . . . . . . . . . p P ◦ / − s S /
28 8.2 0 . ± . − . ± .
09 IP1.14100 Na I 3 p P ◦ / − s S /
68 7.2 0 . ± . − . ± . b P − z D ◦ . . . . . . . . . SB, IP1.15996 Fe I a P − z D ◦
63 8.9 0 . ± . − . ± .
04 FB1.16138 Fe I a P − z D ◦
68 7.9 1 . ± . − . ± .
04 FB1.16348 C I 3 p D − d D ◦ . . . . . . . . . SB1.16442 Fe I a P − z D ◦
59 9.3 0 . ± . − . ± .
04 FB1.16643 C I 3 p S − d P ◦ . . . . . . . . . SB1.16959 Fe I a P − z D ◦
59 7.8 0 . ± . − . ± .
04 Blended withK I line?1.17586 C I 3 p D − d F ◦
41 11.2 3 . ± . − . ± .
03 Blend ofC I transitions1.17891 Fe I b P − z D ◦
54 7.8 1 . ± . − . ± . . . . p P ◦ − s S . . . . . . . . . SB1.18440 Ca II 5 s S / − p P ◦ / . . . . . . . . . SB1.18893 Fe I a P − z D ◦ . . . . . . . . . SB withFe I and C I lines?1.19001 C I 3 p D − d F ◦ . . . . . . . . . SB, Blend ofC I transitions1.19560 Ca II 5 s S / − p P ◦ /
63 9.9 1 . ± . − . ± .
04 FB1.19745 TR . . . . . . . . . . . . . . . . . . b P − z D ◦
30 8.5 . . . . . .
IP, blendedwith TR1.19897 Si I 4 s P ◦ − p D . . . . . . . . . SB, IP1.19972 Si I 4 s P ◦ − p D . . . . . . . . . SB, IP1.20377 Si I 4 s P ◦ − p D
61 9 1 . ± . − . ± . . . . d D − f F ◦
69 9 0 . ± . − . ± . . . . s P ◦ − p D
80 12.2 0 . ± . − . ± . . . . s P ◦ − p D
56 10.4 0 . ± . − . ± .
05 CP, blended ?1.24399 UF . . . . . . . . . . . . . . .
SB1.24665 UF . . . . . . . . . . . . . . .
SB1.25681 C I 3 p P − d P ◦ . . . . . . . . . SB1.25681 [Fe II] a D / − a D / -62 . . . . . . . . . CP, SB1.26200 C I 3 p P − d P ◦ . . . . . . . . . IP1.27870 TR . . . . . . . . . . . . . . .
Weak TR1.28239 H I (Pa β ) 3 − . ± . − . ± .
14 Componentat -649 km.s − . . . . . . . . . . . . . . . . . . S / − P ◦ /
67 8.2 0 . ± . − . ± . . . . S / − P ◦ /
78 6.6 0 . ± . − . ± . . . . b G − z F (cid:12)
47 10.8 0 . ± . − . ± . . . . . . . . . . . . . . . . . . . . . . d D − f F ◦ . . . . . . . . . IB, IP1.50327 Mg I 4 s S − p P ◦ . . . . . . . . . SB1.50484 Mg I 4 s S − p P ◦ . . . . . . . . . SB1.50484 H I (Br 23) 4 − . . . . . . . . . . . . ND1.50436 TR . . . . . . . . . . . . . . . . . . −
22 47 9.9 1 . ± . − . ± .
04 Overlaping TR1.51405 H I (Br 21) 4 −
21 51 14.6 1 . ± . − . ± . . . . −
20 55 19.4 1 . ± . − . ± . . . . −
19 62 17.5 1 . ± . − . ± .
07 FB1.53034 Fe I e D − n D ◦ . . . . . . . . . SB1.53495 [Fe II] a F / − a D / . . . . . . . . . . . . ND 17. Bonnefoy et al.: Multi-epoch images and spectra of the components of Z CMa in the near-infrared during the 2008 outburst
Table 5.
Continuation of Table 4. λ obs Element Transition ∆ V FWHM Flux EW Remarks( µ m) (km . s − ) (Å) (10 − erg . s − . cm − ) (Å)1.53495 H I (Br 18) 4 −
18 57 16.7 1 . ± . − . ± . . . . −
17 64 14 1 . ± . − . ± . . . . −
16 52 11.5 3 . ± . − . ± . . . . CO? 3–0 of X Σ + − X Σ + . . . . . . . . . . . . SB1.57083 H I (Br 15) 4 −
15 52 13.6 . . . . . .
Overlaping TR1.570-1.574 TR . . . . . . . . . . . . . . . . . . p P ◦ − d D . . . . . . . . . SB, IP1.57746 Mg I 4 p P ◦ − d D . . . . . . . . . SB, IP1.57746 CO? 4–1 of X Σ + − X Σ + . . . . . . . . . . . . SB1.58927 H I (Br 14) 4 −
14 136 . . . . . . . . .
Blended withSi I line & TR1.59677 Si I 4 p D − s P ◦ . . . . . . . . . SB1.59871 CO? 5–2 of X Σ + − X Σ + . . . . . . ?? ?? . . . −
13 55 18.8 1 . ± . − . ± .
05 Close TR?1.61962 CO? 6–3 of X Σ + − X Σ + . . . . . . . ± . − . ± . . . . −
12 74 18.8 . . . . . .
CP, SB1.64117 [Fe II] a F / − a D / . . . . . . . . . . . . ND1.68159 H I (Br 11) 4 −
11 75 17.8 2 . ± . − . ± . . . . . . . . . . . . . . . . . . . SB1.68990 C I 3 p D − d F ◦ . . . . . . . . . SB1.70105 He I 3 p P ◦ − d D
48 9.1 . . . . . .
Overlaping TR1.71171 Mg I 4 s S − p P ◦
54 8.8 0 . ± . − . ± . . . . . . . . . . . . . . . . . . . . . . −
10 66 15.6 . . . . . .
Close TR1.747-1.800 TR . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . γ ) 4 − . ± . − . ± . . . . s S / − p P ◦ / . . . . . . . . . IB, IP2.20770 TR . . . . . . . . . . . . . . . . . . s S / − p P ◦ / . . . . . . . . . IB, IP2.29512 CO 2–0 of X Σ + − X Σ + . . . . . . . ± . − . ± . . . . CO 3–1 of X Σ + − X Σ + . . . . . . . ± . − . ± . . . . CO 4–2 of X Σ + − X Σ + . . . . . . . ± . − . ± . . . . . . . . . . . . . . . . . . . . . . CO 5–3 of X Σ + − X Σ + . . . . . . . . . . . . Overlaping TRNotes:
T R : Telluric residuals. FB : Line foot blended with other lines. S B : Line strongly blended. CP : Complex profile. IB : Identified only inemission in the spectrum not corrected for telluric absorptions. IP : Poor accuracy on the estimation of the line position. UF : Unidentified feature. ND Not seen directly.
Table 6.
Variability of the equivalent widths of the Paschen β and Brackett γ lines, and of the ν = −→ CO overtone for theHBe component
Date Instrument EW Pa β EW Br γ EW CO(DD / MM / YYYY) (Å) (Å) (Å)17 / / − . ± . / / − . ± .
10 . . .15 / / − . ± .
08 . . .16 / / − . ± .
08 . . .06 / / − . ± . − . ± . − . ± . / / − . ± . − . ± . − . ± . / / − . ± ..