Limiting Magnitude: 17 (Texp~ 1H, RV ~ 100M/S)

The Right Instrument for your Science Case Henri Boffin ESO

Karen Humby, The Messenger

An amazing suite

Imagers EFOSC2 visible SOFI IR Spectrographs: Resolving Power versus l FORS2 visible HAWK-I NB, JHKs ESPRESSO VISIR M 100 000 NACO AO JHK, Lp UVES CRIRES SPHERE AO visible, JHK

HR HR Multiplex 10 000 FLAMES X-shooter FORS2 MXU LR KMOS IFU (24 arms) VISIR

Resolving power Resolving IR MUSE IFU 1’x1’ MUSE KMOS & SINFONI SINFONI IFU 1000 FLAMES 130 fibers FORS2 & SPHERE VIMOS LR Polarimeters FORS2 100 100 1000 10 000 EFOSC2 Wavelength (nm) SOFI Interferometry SPHERE PIONIER NACO GRAVITY VIRCAM

Attached to VISTA, a 4m telescope Wide-field infrared imager 0.9–2.5 μm Pixel: 0.34” BB and NB filters

FoV: 1. 5 deg2 in 6 pawprints 16 CCDs OMEGACAM

Attached to the 2.6-metre VLT Survey Telescope (VST)

Visible: 350 to 1000 nm

BB and NB filters

32 individual CCDs

Pixel size: 0.21”

FoV: 1 x 1 degree Orion in full VISTA

Optical Infrared FORS2 FORS2

A wonderful German understatement!

It should be called

FOcal Reducer Fast Imager and low dispersion Single and Multi-Object Spectropolarimeter FORFISMOS The Workhorse instrument in Paranal! FORS Iconic Images

FOV: 7’x7’ Pixel scale: 0.125” (2[\!<(28:!0,!Imaging 789:"] FH/!15,/I+!5.345)4!J234K65.5+J!5*!+H/!L,031!L3)1!.34)5+-1/! I36I-63+/1!M0,!3!N05)+!*0-,I/!0M!G/,0!I060-,!OP%"?@A!B>"#@C!<>"#@D!9>"C@C!(>"?@E

789:"!2(F =>"&@?!B>"#@'!<>"#@%!9>"C@#!(>"?@#!G>"&@#

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Remember: FORS2 has a LADC and narrow band filters !"#$%&$'" ! IMAGING

Halley V = 28.2

FORS1+ FORS2+ VIMOS

3 nights

81 images

32284 s Long-slit spectroscopy

330 – 1100 nm

Slits: 0.3” to 2.5”

R ~ 150 – 2600

Mag limit ~ 23 – 24

Fleming 1 PN: host of a post-CE binary

Boffin+ 12, Science MXU

Up to ~100 slits Slit width: 0.1” à30” Slit length < 30” slit shapes: rectangular, circular, and curved

Wavelength range depends on position on CCD (x-axis) Sterzik+ 12 Transmission Spectroscopy

Transiting exoplanets

Tens of ppm precision

Use MXU and measure transit depths as function of wavelength

Sedaghati+ 17 The little but older brother: EFOSC2

ESO Faint Object Spectrograph and Camera 2

FOV: 4.1’x4.1’ Pixel scale: 0.12" The little but older brother: EFOSC2

Arp 271 Imaging (BB, NB)

Hen 2-29

PN M3-4 NGC 3132 The little but older brother: EFOSC2

Arp 271 Imaging (BB, NB) Spectroscopy MXU Polarimetry

Spectrum of a Young Star in M17 The little but older brother: EFOSC2

SN 2018aca Type IA SoFI: Son of ISAAC

0.9-2.5 µm range

Imaging (0.28”/pix) - BB and NB filters.

Spectroscopy: low resolution (R=600) and medium resolution (R=1200-1500), with fixed width slits of 0.6“, 1“ and 2”.

Kilonova AT2017gfo Smartt+ 2017 SoFI: Son of ISAAC

0.9-2.5 µm range

Imaging (0.28”/pix) BB and NB filters.

Spectroscopy: low resolution (R=600) and medium resolution (R=1200-1500), with fixed width slits of 0.6“, 1“ and 2”.

Imaging polarimetry

High time-resolution imaging in burst and fast photometry mode with integration times of the order of a few tens of milliseconds via hardware windowing of the detector array. Ivanov+ 11 HARPS: The exoplanet hunter!

HARPS is an exceptionally STABLE spectrograph 380 to 690 nm Resolving power: 115,000

• RV accuracy: ~ 0.5 m/s on the short term (instrument, guiding, ThAr, atmosphere); ~ 1 m/s on the medium term (2 years);

• Limiting magnitude: 17 (Texp~ 1h, RV ~ 100m/s). HD 10380 C. Lovis et al.: The HARPS search for southern extra-solar planets 13

12 C. Lovis et al.: The HARPS search for southern extra-solar planets 4

0.6 5 to 7 planets in one 0.5 (a) 3 0.4 0.3 system: 0.2 2 0.1 0 0.3 1 0.25 (b) • (1 Super-Earth, with 1.4 0.2 0.15 ecc 0 0.1 Earth masses, orbital 0.05 y (AU) 0 0.3 -1 period of 1.18 days) 0.25 (c) 0.2 0.15 -2 • 5 Neptune-like planets, 0.1 0.05 0 -3 with masses between 0 2 4 6 8 10 12 14 16 18 20 time (kyr) -4 13 and 25 Earth Fig. 8. Evolution of the eccentricities-4 of the planets-3 b (red)-2 and c (green)-1 0 1 2 3 4 during 20 kyr for three different models. In the top picture the initial eccentricity of planet b is set at zero, but mutual gravitational perturba- x (AU) masses, orbital periods tions increase its value to 0.4 in less than 2 kyr (Table 3). In the middle figure we included general relativity, which calms down the eccentricity variations of the innermost planet, but still did not preventtheeccentric- ity of planet d toFig. reach high 11. values.Long-term In the bottom figure evolution we use a model of the HD 10180 planetary system over from 6 to 600 days where the eccentricities of both planets were previously damped by tidal dissipation (Table10 6). ThisMyr last starting solution is stable with at least thefor orbital 10 Myr. solution from Table 6. The panel shows Fig. 13. The 15 planetary systems with at least three known planets as • 1 Saturn-like planet, aface-onviewofthesystem.x and y are spatial coordinates in a frame of May 2010. The numbers give the minimal distance between adjacent 1 centered on the star. Present orbital solutions are traced with solid lines planets expressed in mutual Hill radii. Planet sizes are proportional to with 65 Earth masses, 0.8 and each dot corresponds to the position of the planet every 100 kyr. log (m sin i). 0.6 The semi-major axes (in AU) are almost constant, but the eccentricities undergo significant variations (Table 7). amp orbital period of 2200 0.4 (type II) through angular momentum exchange with the gaseous 0.2 disk; disk evolution and dissipation within a few Myr; dynamical days 0 0 0.2 0.4 0.6 0.8 1 interactions between protoplanets leading to eccentricitypump- time (Gyr) ing, collisions or ejections from the system. It is extremelychal- Fig. 9. Tidal evolution of the amplitude of the proper modes u1 (red), u2 lenging to build models that include all these effects in a con- (green), u3 (blue), and u4 (pink) resulting from the tidal dissipation on planet b with k2/Q = 0.0015. sistent manner, given the complicated physics involved and the Hara+ 16 à 6 planets secure scarce observational constraints available on the early stages of above do probably not reach down to the Neptune-mass range planet formation. In particular, attempts to simultaneously track yet, preventing us from having a sufficiently complete picture the formation, migration and mutual interactions of several pro- Lovis+ 2011 of them. Nevertheless, considering the well-observed casesfol- lowed at the highest precision, the RV technique shows here its toplanets are still in their infancy. Observational resultsonthe ability to study the structure of planetary systems, from gasgi- ants to telluric planets. global architecture of planetary systems may therefore provide important clues to determine the relative impact of each process.

Fig. 7. Phased RV curves for all signals in the 7-Keplerian model. In From the observational point of view, the least that can be each case, the contribution of the other 6 signals has been subtracted. said is that planetary systems display a huge diversity in their properties, which after all is not surprising given their complex formation processes. Fig. 13 shows planet semi-major axes ona logarithmic scale for the 15 systems known to harbour at least three planets. Systems are shown in increasing order of their mean planetary mass, while individual planet masses are illus- Fig. 12. Possible location of an additional eighth planet in the HD 10180 trated by varying dot size. The full range of distances covered system. The stability of a small-mass particle in the system is analyzed, by each planet between periastron and apastron is denoted by 1 for various semi-major axes and eccentricities, and for K = 0.78 m s− . ablackline.Theminimaldistancebetweeneachneighbouring The stable zones where additional planets could be found are the “dark pair of planets is also given, expressed in units of the mutualHill blue” regions. radius which is deﬁned as:

1/3 a1 + a2 m1 + m2 R , = . (10) H M 2 ! 3 M " The dynamical architecture of planetary systems is likely to ∗ convey extremely useful information on their origins. The dom- As shown by Chamberset al. (1996),the instability timescale inant planet formation scenario presently includes severalphys- of a coplanar, low-eccentricity multi-planet system is related to ical processes that occur on similar timescales in protoplane- the distance between planets, expressed in mutual Hill radii, in tary disks: formation of cores, preferentially beyond the ice line, arelativelysimplemanner,andapproximate’lifeexpectancies’ through accretion of rocky and icy material; runaway gas accre- of planetary systems can be estimated based on these separa- tion on cores having reached a critical mass, rapidly forminggi- tions, the number of planets, and their masses. The Chambers ant planets; inward migration of cores (type I) and giant planets et al. simulations do not extend to masses above 3 M or to ∼ ⊕ Proxima b

K ≈ 1.4 m/s

M sin i ≈ 1.3 ME P = 11.2 d

Anglada-Escudé + 16 Big Brother: ESPRESSO

380-686 nm

R= 120,000 – up to 220,000 (1 UT) or 30,000 (4UTs)

RV precision < 10 cm/s (1 UT) < 1-5 m/s (4 UTs)

S/N=260 in one hour on a V=12 star (1 UT) To be offered soon!

S/N = 20 on a V = 20 object in one hour (4 UTs) UVES

Echelle Spectrograph

300–500 nm (blue arm) 420–1100 nm (red arm)

Spectral resolution ~ 40,000 up to 80 000 (blue arm) – 110 000 (red arm)

SN 1987A in 1999 Rapid Response Mode

Observe within minutes of discovery

GRB 060418 Vreeswijk+ 107

Also for kilonovae, GW Detecting CO molecules

8h spectrum of a distant quasar

Detect CO molecules 11 billion ly away

Measure Temperature of Universe at this epoch: 9.15 K (in excellent

Srianand+ 08 agreement with Big Bang) A Pristine Star

Extremely metal-poor star ~13 billion year old Caffau+ 11 X-SHOOTER

Echelle spectrograph

300–2500 nm

R ~ 4000–17 000 (various slit widths) Mag limit ~ 21

Three arms: UVB VIS Near-IR X-SHOOTER

Colliding Neutron Stars

ESO/E. Pian et al./S. Smartt & ePESSTO FLAMES: multi-fibre robot

370–950 nm

R=7000–24 000 with GIRAFFE (up to 130 fibres)

R=47 000 with UVES (up to 8 objects)

FoV: 25’ diam.

Fibre fed: 1”or 1.2” diam IFU also possible (2” x 3”) FLAMES

NGC 6752 Second generation of stars fail to enter AGB phase

Campbell+ 13 MUSE

Integral Field Spectrograph 24 spectrographs

24 x 48 = 1152 mini-slits

365–930 nm

R ~ 4000

Pixel scale: 0.2” FoV: 1’ x 1’ With or without AO (AOF) MUSE MUSE: Redshift 5 Galaxy

Pillars of Creation (M16)

McLeod+ 15 3D Structure Lifetime: 3 Myrs A.Cluster Mehner et al.: A jet emergingSH from2-266 the evolved B[e] star MWC 137

A. Mehner et al.: A jet emerging from the evolved B[e] star MWC 137

30 16 NaI CaII CaII CaII 24: F6−8V 23 25 14 25: F6−8V

25 HeI HeI HeII HeI HeI HeI HeI OI 26: F1−3V 1: sgB[e] 15.295 29 27: F1−3V 2: B2−4V 12 3: B3−5V 28: F6−8V 20 4: B8−9V 5: B6−8V 10 29: F1−3V 20 6: B7−8V 30 7: B8−9V 30: F1−3V 8: B7−8V 9: B8−9V 8 15 31: F2−5V 10: B8−9V 32: F2−3V 37 11: A1−3V 12: B8−9V 6 33: F7−8V 13: B8−9V Relative Flux + Const. Relative Flux + Const. 10 14: A3−4V 34: G5V 12 35 15: B9−A0V 16: B9−A1V 4 35: G5V 3

17: A0−4V ) 31 15 18: A0−4V

° 33 36: G5V 5 19: A7−9V 2 37: G5V 20: A1−2V 15.285 38: G5V 21: A1−2V 22: A1−3V 0 39: G5V 19 0 23: A9V−F0V 22 4 5000 6000 7000 8000 9000 5000 6000 7000 8000 9000 13 14 Wavelength (Å) Wavelength (Å) 5 11 1 8 Fig. 4. Spectra of the brightest cluster stars extracted from the MUSE data cube. The spectrum #1 is of9 MWC 137. Spectra #2 to #23 are B- 16 10 and A-type stars in the cluster. Their spectral classification, indicated next to each spectrum, was based primarily on the H↵,H, and O I 7774 equivalent widths. Spectra #24 to #39 are F- and G-type stars. Their classification is based primarily on the Ca II 8542 and Na I 5889 equivalent 18 6 widths. The numbering is consistent with Table 2 and Figure 1. Grey areas block regions of the spectra24 with strong residuals from the telluric bands correction. 21 7 2 1 tion. We determine a systemic cluster velocity of 50 20 km s For the hot stars (spectral types B and A), we based our clas- ± 34 based on the radial velocitiesDeclination ( of the Paschen lines (hot stars) sification mostly on H↵ and H equivalent widths. The strength 27 and Ca II and Ca I lines (cool stars). The spectra have only of O I 7774 was used36 to confirm the spectral subclass. The clus- few spectral lines and there are residuals from the nebula ter does not contain any O-type stars. Two objects show weak 32 emission for the broad hydrogen lines. This results in less ac- He I absorption and we determined their spectral type to early curate radial velocity measurements than expected for the B. For the cooler stars (spectral types F and later), we based our spectral resolutionFor of MUSE. 15.275each star classificationin the on the equivalent cluster, widths of Ca II 8542 and Na I We determined the stellar spectral types with the help of the 5889. The second to last column in table Table 2 lists our spec- ESO UVES Paranal Observatory Project (POP) library of high- tral classifications. resolution spectra (R 80 000, Bagnulo et al. 2003) and the In addition to MWC 137, we found two stars with potential Indo-US Librarywe of low-resolution⇠ have spectra (R a5 000, spectrum Valdes H↵ emission (stars #4 and #16à in Table 2). H↵ and O I 8446 et al. 2004). The spectral resolving power of a range⇠ of template emission remain after the nebular subtraction, while other neb- 39 spectra, covering late-O to mid-G spectral subclasses, was de- ular lines are reasonable well removed. This may indicate cir- creased to R 3 000 to match our spectra. Equivalent widths of cumstellar material. The hydrogen Paschen lines are very broad 17 lines used forspectral⇠ the spectral classification (see below)classification were mea- in absorption and suggest a high valueà of log g, indicating a lu- sured in these template spectra and complemented by line equiv- minosity class V and main sequence evolutionary phase. These 26 alent widths listed in Jaschek & Jaschek (1987) and Didelon two stars may be Be stars. However, the nebula emission varies (1982). MUSE does not cover the critical wavelength regions strongly on small spatial scales and bad nebular subtraction may bluewards ofextinction 4750 Å commonly used for stellar classification à of magnitudeleave remnants that resemble narrow H ↵ emission.à Follow-up hot stars (Walborn & Fitzpatrick 1990). The spectral range and observations with higher spectral resolution are needed to deter- 28 resolving power do also not permit the classification of G-type mine whether these stars are indeed Be stars. star subclasses. Our spectral classification is thus only accurate The right panel of Figure 1 indicates the location of the clus- to within a fewisochrones spectral subclasses (Table 2). A review with à spec- terage stars on sky. It demonstratesof cluster that the higher mass stars are tra at bluer wavelengths is needed. located preferentially towards the cluster center, while the lower 38 4 15.265 94.705 94.695 94.685 94.675 Right Ascension (°) Mehner+ 15

Fig. 1. Left: Composite VRI color image (1.70 1.70) of the five MUSE fields covering the cluster SH 2-266, the central star MWC 137, and its circumstellar nebula. Right: Spatial distribution⇥ of the brightest cluster stars (blue diamond in the center: MWC 137; black squares: A- and B-type stars; red squares: F- and G-type stars; black open squares: unclassified stars; arrow: jet extension, see Section 3.4). All stars for which a spectral type was obtained are indicated with a number, consistent with Table 2.

MWC 137 is most likely in an extremely short-lived phase evolv- at each position of the cluster SH 2-266 and the large-scale neb- ing from a B[e] supergiant to a blue supergiant with a bipolar ula around MWC 137. ring nebula. However, the 12C/13C ratio is still compatible with The data were reduced using version 1.1.0 of the MUSE that of a protoplanetary nebula and a main sequence star with a standard pipeline (Weilbacher et al. 2012). Bias, arc, and flat- mass M > 9 M , because the 12C/13C ratio drops for such stars field master calibrations were created using a set of standard already on the main sequence (Kraus 2009). The observed value calibration exposures. Bias images were subtracted from each of the 12C/13C ratio is indicative, but by itself it does not provide science frame. The science frames were flat-fielded using the strong evidence for the object’s evolutionary stage. master flat-field and renormalized using a flat-field taken right We investigate the cluster SH 2-266 to establish MWC 137’s after the science to account for the temperature variations in the membership and constrain the star’s evolutionary stage using illumination pattern of the slices. An additional flat-field correc- data obtained with the new Very Large Telescope Multi Unit tion was performed using the twilight sky exposures to correct Spectroscopic Explorer (VLT MUSE; Bacon et al. 2010). While for the di↵erence between sky and calibration unit illumination. MWC 137 is generally associated with this cluster (Testi et al. A geometrical calibration and the wavelength calibration solu- 1999), no detailed study of the cluster properties exist. In Section tion were used to transform the detector coordinate positions to 2 we describe the observations. In Section 3 we present our re- wavelengths and spatial coordinates. The astrometric solution, sults regarding the cluster and MWC 137. We analyze a col- flux calibration, and telluric correction were then applied. Sky limated outflow, which we discovered in the MUSE data and subtraction was performed only in the post-reduction analysis of which probably originates from MWC 137. In Section 4 we dis- the individual stellar spectra. A data cube was produced from cuss the potential origin of this jet with respect to the evolution- each pixel-table using a three-dimensional drizzle interpolation ary phase of MWC 137. Section 5 gives a short conclusion. process, which also rejects cosmic rays. All data cubes were sampled to 0. 2 0. 2 1.25 Å. 00 ⇥ 00 ⇥ 2. Data and analysis The MUSE standard pipeline also delivers field-of-view images in di↵erent filters when the filter transmission curves (in The cluster SH 2-266 with the B[e] star MWC 137 in the cen- our case Johnson V, Cousins R, and Cousins I) are provided, ter was observed with VLT MUSE on 2014-11-07 and 2014-12- see Figure 1 for a composite VRI image. We used the Aperture 01. MUSE is a spectrograph composed of 24 integral field units Photometry Tool v.2.4.72 to determine the VRI magnitudes of (IFUs) that sample a 10 10 field of view (Bacon et al. 2010). The the brightest cluster members. Our MUSE data do not provide ⇥ instrument covers almost the full optical domain (4800–9300 Å) absolute magnitudes and we thus used V = 11.95 mag for with a spectral resolving power of R 3 000. The spatial reso- MWC 137 as a reference (Reed 2003). We found no I-band mag- ⇠ lution for MUSE’s currently o↵ered wide field mode is 000. 3–000. 4 nitude reported in the literature. We therefore compare the V–I (the pixel scale is 000. 2 per pixel). We obtained seeing-limited colors of the stars for which we determined the spectral type and spatial resolutions between 000. 6 and 000. 7. extinction (Section 3.2) to the V–I colors reported in Pecaut & Three exposures of 2 s, 20 s, and 300 s were obtained with Mamajek (2013) to ensure our values are consistent with litera- MWC 137 in the center of the field of view of MUSE. Additional ture values. QFitsView3 was used to extract the stellar spectra us- 600 s exposures were acquired with o↵sets of 3500 in each of ing circular apertures from which an annulus mean sky/nebular the cardinal directions (see Figure 1). The limiting point source magnitude in our data is V 24 mag. MUSE provides the first 2 http://www.aperturephotometry.org/aptool/ ⇡ 3 complete and accurate set of optical line ratios and line centroids http://www.mpe.mpg.de/õtt/dpuser/qfitsview.html

2 A.Cluster Mehner et al.: A jet emergingSH from2-266 the evolved B[e] star MWC 137

A. Mehner et al.: A jet emerging from the evolved B[e] star MWC 137

30 16 NaI CaII CaII CaII 24: F6−8V 23 25 14 25: F6−8V

17: A0−4V ) 31 15 18: A0−4V

° 33 36: G5V 5 19: A7−9V 2 37: G5V 20: A1−2V 15.285 38: G5V 21: A1−2V 22: A1−3V 0 39: G5V 19 0 23: A9V−F0V 22 4 5000 6000 7000 8000 9000 A. 5000 Mehner 6000 et al.: 7000 A jet emerging8000 9000 from the evolved B[e] star13 MWC 137 14 Wavelength (Å) Wavelength (Å) 5 11 1 8 Fig. 4. Spectra of the brightest cluster stars extracted from the MUSE data cube. The spectrum #1 is of9 MWC 137. Spectra #2 to #23 are B- 16 10 and− A-type10 stars in the cluster. Their spectral classification, indicated next to each spectrum, was based primarily on the H↵,H, and O I 7774 equivalent widths. Spectra #24 to #39 are F- and G-type stars. Their classification is based primarily on the Ca II 8542 and Na I 5889 equivalent 18 6 widths. The numbering is consistent with Table 2 and Figure 1. Grey areas block regions of the spectra24 with strong residuals from the telluric UKIDSS bands correction. 1 Myr 21 −8 7 2 1 20 M tion. We determine a systemic cluster velocity of 50 20 km s For the hot starssun (spectral types B and A), we based our clas- 1 ± 34 based on the radial velocitiesDeclination ( of the Paschen lines (hot stars) sification mostly on H↵ and10 H equivalent Myr widths. The strength 27 and Ca II and Ca I lines (cool stars). The spectraMWC137 have only of O I 7774 was used36 to confirm the spectral subclass. The clus- −6 15 Msun few spectral lines and there are residuals from the nebula ter does not contain any O-type stars. Two objects show weak 32 5 Myr emission for the broad hydrogen lines. This results in less ac- He I absorption and we determined their15 spectralMyr type to early curate radial velocity measurements than expected for the B. For the cooler stars (spectral types F and later), we based our spectral−4 resolution of MUSE. classification on the equivalent widths of Ca II 8542 and Na I 15.275 10 Msun We determined the stellar spectral types with the help of the 5889. The second to last column in table Table 2 lists our spec- ESO UVES Paranal Observatory Project (POP) library of high- tral classifications. resolution spectra (R 80 000, Bagnulo et al. 2003) and the In addition to MWC 137, we found two stars with potential Indo-US−2 Library of low-resolution⇠ spectra (R 5 000, Valdes H↵ emission (stars #4 and #16 in Table 2). H↵ and O I 8446 0.5 et al. 2004). The spectral resolving power of a range⇠ of template emission remain after the nebular subtraction, while other neb- 39 H spectra,V covering late-O to mid-G spectral subclasses, was de- ular lines are reasonable well removed. This may indicate cir- − M creased 0 to R 3 000 to match our spectra. Equivalent widths of cumstellar material. The hydrogen Paschen lines are very broad J 17 MWC lines used for⇠ the spectral classification (see below) were mea- in absorption and suggest a high value of log g, indicating a lu- 16 sured in these template spectra and complemented by line equiv- minosity class V and main sequence evolutionary phase. These 26 alent widths listed in Jaschek & Jaschek (1987) and Didelon two stars may be Be stars. However, the nebula emission varies (1982). 2 MUSE does not cover the critical wavelength regions strongly on small spatial scales and bad nebular subtraction may 8 leave remnants that resemble narrow H↵ emission. Follow-up bluewards of 4750 Å commonly used for stellar classification of 4 hot stars (Walborn & Fitzpatrick 1990). The spectral range and observations with higher spectral resolution are needed to deter- 0 28 6 resolving 4 power do also not permit the classification of G-type mine whether these stars(pre are− indeedmain Be stars.sequence) star subclasses. Our spectral classification is thus only accurate The right panel of Figure 1 indicates the location of the clus- to within a few spectral subclasses (Table 2). A review with spec- ter stars on sky. It demonstrates that the higher mass stars are tra at bluer wavelengths is needed. located preferentially towards the cluster center, while the lower D = 5.2 kpc 6 1 Myr 38 4 15.265 3 Myr old 8 10 Myr 94.705 94.695−0.5 94.685 94.675 −0.5 0 0.5 1 1.5 2 2.5 3 −0.5 0 0.5 1 V−I Right Ascension (°) H−K Mehner+ 15 Fig. 6. MV versus V–I color-magnitude diagram of the brightest SH 2- Fig. 7. J–H versus H–K color-color diagram from the 266 cluster stars derived from the MUSE data. Black squares are B- and UKIDSSDR8PLUS database. Black squares are B- and A-type Fig. 1. Left: Composite VRI color image (1.70 1.70) of the five MUSEA-type stars, fields red squares covering are cooler F- the and G-type cluster stars, corrected SH for2-266,stars, red the squares central are cooler F- star and G-type MWC stars. The 137, open gray and squares its ⇥ their individual extinction. The arrow at MWC 137 indicates the po- are unclassified stars. The theoretical stellar evolution isochrone for circumstellar nebula. Right: Spatial distribution of the brightest clustersitions stars in the color-magnitude(blue diamond diagram from in the the average center: inner cluster MWC5 Myr from 137; the Geneva black stellar squares: models is shown A- (Ekstr andom¨ et B-type al. 2012, estimate of E(B–V) = 1.2 mag to the estimate for MWC 137 of E(B– Yusof et al. 2013). MWC 137, #10, #23, #25, and #28 are saturated stars; red squares: F- and G-type stars; black open squares: unclassifiedV) 1 stars;.8 mag. The arrow: open gray squares jet extension, are unclassified stars see corrected Sectionin the UKIDSS 3.4). images All and stars their JHK forvalues which were obtained a spectral from the for a⇠ standard extinction of E(B–V) = 1.2 mag. Theoretical stellar evo- 2MASS point source catalogue (Cutri et al. 2003; the 2MASS cata- type was obtained are indicated with a number, consistent with Tablelution 2. isochrones for 1 Myr, 10 Myr, and 15 Myr (gray dashed curves) logue does not contain the star #10, which is thus not displayed here). and the evolutionary paths for 10 M , 15 M , and 20 M stars from The data were dereddened adopting the E(B–V) values determined the Geneva stellar models are shown (Ekstrom¨ et al. 2012, Yusof et al. in Section 3.2 and the relations E(J–H)/E(H–K) = 2.09 mag (Jones 2013). 1 Myr and 10 Myr pre-main sequence theoretical isochrones for & Hyland 1980) and E(J–H)/E(B–V) = 0.33 mag (Lee 1970). For star masses of 0.1–7 M from Siess et al. (2000) are also shown (gray MWC 137 the dereddening vector is shown from E(B–V) = 1.2 mag to dotted curve). 1.8 mag. For the unclassified stars a default value of E(B–V) = 1.2 mag MWC 137 is most likely in an extremely short-lived phase evolv- at each position of the clusterwas SH assumed. 2-266 and the large-scale neb- ing from a B[e] supergiant to a blue supergiant with a bipolarare observedula to aroundhave excess near-infrared MWC emission 137. (compare to Koornneef 1983) possibly due to hot dust and in the case 12 13 main sequence cut-o↵ is around 8 M . This is also in agree- ring nebula. However, the C/ C ratio is still compatible withof MWC 137 a disk-like structure may also contribute. Near- The data were reduced usingment with version the near-infrared 1.1.0 colors of of the cluster the stars. MUSE infrared emission is present for the stars #4 and #16, confirming 3. MWC 137 is not following the evolution of a single star. that of a protoplanetary nebula and a main sequence star withthe a presence of circumstellar material, which the H↵ emission standard pipeline (WeilbacherBased et al. on the 2012). fact that no stars Bias, with protoplanetary arc, and disks flat- are in their spectra suggested (Section 3.2). For the two other stars 12 13 found in this cluster and the near-infrared colors, the cluster mass M > 9 M , because the C/ C ratio drops for such starswith near-infraredfield excess, master stars #6 andcalibrations #8, there is no evidence were created using a set of standard is older than 3–5 Myrs. of optical line emission. 4. The cluster may have had several star-formation episodes. In already on the main sequence (Kraus 2009). The observed valueFigurecalibration 6 indicates that a few exposures. F- and G-type stars Bias may not images were subtracted from each this case a single age determination is meaningless. 12 13 have reached the main sequence yet, but have (pre-main se- of the C/ C ratio is indicative, but by itself it does not providequence) agesscience between 1 Myr frame. and 10 Myr The based on science the tracks by frames were flat-fielded using the Siess et al. (2000). These stars (#24,25,30,31 at MV 0.5 mag) 3.4. The jet strong evidence for the object’s evolutionary stage. are locatedmaster along the north-east flat-field rim of the and nebula, renormalized⇠ implying using a flat-field taken right a later star-formation episode. A shock front originating from We discovered a collimated outflow in the MUSE data that ap- We investigate the cluster SH 2-266 to establish MWC 137’sMWC 137after could have the caused science a second episode to of account star-formation forpears the to originate temperature from MWC 137 at variations a position angle of in 18–20 the membership and constrain the star’s evolutionary stage usingin this cluster. However, stars #24 and #25 have a low extinction (Figure 8). We detect the jet in H↵, [N II] 6548,6583, and [S II] comparedillumination to the other two stars and pattern also a slightly of thedi↵erent slices.6716,6731, An but additional not in [O I] 5577,6300. flat-field The extension correc- of the jet data obtained with the new Very Large Telescope Multi Unitradial velocity and thus may lie in front of the cluster. The near- as traced by the [N II] 6583 line is about 6600, which corresponds infrared datation show was no evidence performed for these stars being using pre-main theto twilight a length of 1.7 pc skyprojected exposures on the plane of sky,to if correct we adopt Spectroscopic Explorer (VLT MUSE; Bacon et al. 2010). Whilesequence objects.for the di↵erence between skya distance and of calibration5.2 kpc. The width unit of the jet illumination. is with 3 pixels not To summarize the di↵erent possibilities related to the age of resolved, but corresponds to the seeing of about 0.600. The jet ex- the cluster and its very probable member MWC 137: tends almost across the entire cluster (Figures 1 and 8). The jet MWC 137 is generally associated with this cluster (Testi et al. A geometrical calibration andoutflow, the especially wavelength the northern outflow, calibration shows deviations solu- from 1999), no detailed study of the cluster properties exist. In Section1. MWCtion 137 is following were the used evolution to of a singletransform massive star. thea linear detector trace and at least coordinate six knots can be identified positions (labelled ‘a’ to Its position in the Hertzsprung-Russell (HR) diagram is best to ‘f’ in Figure 8). This may hint at a velocity variability and a 2 we describe the observations. In Section 3 we present our re-fit withwavelengths10–15 Myr isochrones, whichand have spatial a main sequence coordinates.precession of the The jet axis (Raga astrometric et al. 2001). solution, turn-o↵ around 10–15 M . This is in agreement with the The velocity structure within the jet is resolved over about near-infrared colors of the cluster stars. five MUSE channels. The jet material expands with projected ve- sults regarding the cluster and MWC 137. We analyze a col- flux calibration, and telluric correction were1 then applied.1 Sky 2. MWC 137 is not following the evolution of a single star. The locities between 148 km s to 433 km s towards the north 1 1 limated outflow, which we discovered in the MUSE data andHR diagramsubtraction is best fit with a was30 Myr performedisochrone, where the onlyand between in the+250 post-reduction km s to +478 km s towards analysis the south. A of which probably originates from MWC 137. In Section 4 we dis- the individual stellar spectra. A data cube was produced from 6 cuss the potential origin of this jet with respect to the evolution- each pixel-table using a three-dimensional drizzle interpolation ary phase of MWC 137. Section 5 gives a short conclusion. process, which also rejects cosmic rays. All data cubes were sampled to 0. 2 0. 2 1.25 Å. 00 ⇥ 00 ⇥ 2. Data and analysis The MUSE standard pipeline also delivers field-of-view images in di↵erent filters when the filter transmission curves (in The cluster SH 2-266 with the B[e] star MWC 137 in the cen- our case Johnson V, Cousins R, and Cousins I) are provided, ter was observed with VLT MUSE on 2014-11-07 and 2014-12- see Figure 1 for a composite VRI image. We used the Aperture 01. MUSE is a spectrograph composed of 24 integral field units Photometry Tool v.2.4.72 to determine the VRI magnitudes of (IFUs) that sample a 10 10 field of view (Bacon et al. 2010). The the brightest cluster members. Our MUSE data do not provide ⇥ instrument covers almost the full optical domain (4800–9300 Å) absolute magnitudes and we thus used V = 11.95 mag for with a spectral resolving power of R 3 000. The spatial reso- MWC 137 as a reference (Reed 2003). We found no I-band mag- ⇠ lution for MUSE’s currently o↵ered wide field mode is 000. 3–000. 4 nitude reported in the literature. We therefore compare the V–I (the pixel scale is 000. 2 per pixel). We obtained seeing-limited colors of the stars for which we determined the spectral type and spatial resolutions between 000. 6 and 000. 7. extinction (Section 3.2) to the V–I colors reported in Pecaut & Three exposures of 2 s, 20 s, and 300 s were obtained with Mamajek (2013) to ensure our values are consistent with litera- MWC 137 in the center of the field of view of MUSE. Additional ture values. QFitsView3 was used to extract the stellar spectra us- 600 s exposures were acquired with o↵sets of 3500 in each of ing circular apertures from which an annulus mean sky/nebular the cardinal directions (see Figure 1). The limiting point source magnitude in our data is V 24 mag. MUSE provides the first 2 http://www.aperturephotometry.org/aptool/ ⇡ 3 complete and accurate set of optical line ratios and line centroids http://www.mpe.mpg.de/õtt/dpuser/qfitsview.html

2 KMOS

The octopus of Paranal… But with 24 cryogenic arms… and each is an IFU KMOS

0.8µm to 2.5µm IZ, YJ, H, K, H+K

R = 2000–4200

24 IFUs to understand how galaxies grew and in 7.2 arcmin diameter circle evolved in the early cosmos Each IFU: 2.8" x 2.8” 0.2" x 0.2” pixels KMOS3D

MPE 3D HST and CANDELS HAWK-I: High Acuity Wide field K- HAW K-I High Acuity Wide field K-band Imager band Imager

near-infrared 0.85-2.5 µm wide-field imager 7.5’ x 7.5’ Pixel scale 0.1064''

HAWK-I: High Acuity Wide field K- HAW K-I High Acuity Wide field K-band Imager band Imager

near-infrared 0.85-2.5 µm wide-field imager 7.5’ x 7.5’ 1 hour Pixel scale 0.1064'' integration

HAWK-I: High Acuity Wide field K- HAW K-I High Acuity Wide field K-band Imager band Imager

near-infraredY, J, H, K0.85-2.5 and 6 µNBm wide-fieldfilters imageravailable Pixel scale 0.1064''

A gallery of Spiral Galaxies Grosbøl ESO/P. ESO/P. SINFO N I MUSE

RangoRango dede λλ 1.1 – 2.5 µm 0.5 – 0.9 µm

CampoCampo 0.8, 3, 8 arcsec 1 arcmin

AOAO SI NO (por ahora...) AO Facility on UT4 (AOF) FU TU RO

Tarentula without/with AOF

Down to 0.2” resolution à Better resolved à Deeper VISIR: VLT Imager and Spectrometer for mid-Infrared

Mid-Infrared imager, spectrograph

N-band 8–13μm Q-band 16.5–24.5 μm

Pixel scale: 45-76 mas

FoV: 36x38” or 60x60”

Spectroscopy: 350 to 25,000 ‘Hot’ South Pole of Neptune

Troposphere

Stratosphere

T T+6.3h The Fried Egg Nebula

Yellow hypergiant

Huge dusty double shell

ESO/E. Lagadec NAOS-CONICA (NACO)

0.8–5 μm

Adaptive-optics-assisted imaging imaging polarimetry, and coronagraphy

Pixel scale: 13, 27 or 54 mas

Equipped with one infrared (0.8-2.5 µm) and one visual wavefront sensor (0.45-1.0 µm). Reference star can be off-axis!

Can provide Strehl of 50% with V=12 reference star Magnitude limit ~ 23.5-24 in J, H, K Image of an “Exoplanet”

Brown Dwarf “2M1207” 25 Jupiter-masses 8 million years old “Giant Planet Candidate Companion (GPCC)” 100 x fainter 1000 ºC 55 AU distance 5 Jupiter-masses TW Hydrae Association 230 light-years Chauvin et al., 2005 Water molecules Centro GalacticoGalactic Centre

NACO 1’’ ~ 0.04 pc

(0.012 ly) SINFONI

Near-infrared integral field spectrograph with adaptive optics capabilities

1–2.5 μm

R ~ 1500–4000

Pixel scale: 25 to 250 mas

Can use a Laser Guide Star FoV: 0.8”x0.8” to 8”x8”

Mag limit ~ 18-20 Centro GalacticoGalactic Centre

SINFONI SPHERE

High-contrast Imaging (coronagraph), low-resolution spectroscopy (R~30-350), polarimetry

Extreme AO: Strehl ratio of ~75% in H-band

On-axis reference star (R < 13 mag)

In visible and in near-IR: 0.5–2.32 µm

FoV: 1.73"x1.73“ (IFU), 3.5”x3.5” (optical) to 11’’ x 11” (IR) Four planets around HR 8799 Four planets around HR 8799! Protoplanetary discs Better than Hubble

Sphere/IRDIS% VY CMa, a red hyper-giant PIONIER at the VLTI

4 telescopes

Visibilities on 6 baselines and 4 closure phases simultaneously Highly efficient (1 point ~ 5-7 min) A typical calibrated sequence is 45 min PIONIER: Measuring Stars

Studied 6 symbiotic or related stars Visibilities

Boffin+ 2014 Masses of SB2 components 7 the RV of the primary components, for the ancient and for Table 5. The combined VB+SB2 solutions; For consistencyMasses of SB2 components 3 the new measurements. The systemic velocity, V0,isderived with the SB orbits and with the forthcoming astrometric orbit, HIP 14157 HIP 20601 HIP 117186 in the system of the new RV measurements of the primary ω and Ω both2 refer to the motion of the primary component.4 N component. 0 0 2 N

-2 E For HIP 14157 and HIP 20601, F2 = −0.041 and −0.16, HIP 14157-5 HIP 206010 HIP 117186E respectively, indicating that both sets of measurements are Application-4 to Binaries

X (mas) -10 -2 P (days) 43.32032 156.38020N Masses 85.8238 of SB2 components 3 quite compatible. For HIP 117186, F2 =0.78, since some -6 ± 0.00013-15 ± 0.00026 -4 ± 0.0012 -8 HIP 14157 HIP 20601E HIP 117186 discrepancies appear between the SB2 orbit derived from 2 4 T0 (JD-2400000) 51487.5005 56636.6713-6 56402.576 -10 N -20 the new measurements and the preceding one; the most im- 4 2 0 -2 -4 -6 -8 10 5 0 -5 4 2 0 -2 -4 -6 0 ± 0.00790 ± 0.0027 2 ± 0.072 N Y (mas) Y (mas) Y (mas) portant is the mass ratio, which is 0.771 ± 0.021 with the -2 E 0.05 e 0.2 0.75940.2-5 0.851480 0.32702 E 0.1 0.1 (mas) ancient measurements and 0.844 ± 0.012 with our observa- a -40 ± 0.00100 ± 0.00025 0 ± 0.00068 -0.1 -0.1

X (mas) -2 −1 -0.2 -0.2-10 (O-C) N tions. Nevertheless, The SB2 orbit obtained from both sets V0 (km s )30.74341.623-19.89-6 -0.05 0.2 0.2 0.05 -15 -4 0.1-8 ± 0.0910.1 ± 0.014 ± 0.33 of RVs is nearly similar to that derived from our measure- (mas) E b 0 0 0 o -0.1 -0.1 -6 ments alone, but the period is much more accurate, thanks ω1 ( )174.69202.026176.07-0.2-10 -0.2-20 Halbwachs+ 15 (O-C) 4 2 0 -2 -4 -6 -8 10 5 0 -5 -0.05 4 2 0 -2 -4 -6 56950Y 57000(mas) ± 570500.17 56950±Y 57000(mas)0.086 57050 56940± 0.32 56960Y (mas) 56980 57000 to the extension of the timespan covered by the observations. T - 2,400,000 (JD) T - 2,400,000 (JD) T - 2,400,000 (JD) o 0.05 Ω ( ;eq.2000)1500.2 19.1411500.2 340.52640 16.928 1 0.1 0.1

The new SB2 orbits are presented in Table 4. (mas) a 1000 0 200 ± 0.082100 ± 0.058 ± 0.047 -0.1 -0.1 0 -0.2 -0.2 o (O-C) 50 -0.05 i ( )92.24103.13888.05450 0.05-20 0.20 0.2 RV (km/s) 0.1 0.1 -40 (mas) ± 0.18 0 ± 0.077 ± 0.043 b -500 0 0 a -0.1 -0.1 -60 -0.2 -0.2

a (mas)(O-C) 5.810 11.339 4.677 4 THE MASSES 0 0.5 1 0 0.5 1 -0.05 0 0.5 1 56950Phase 57000± 57050 56950±Phase 57000 57050 ±56940 56960Phase 56980 57000 T - 2,400,000 (JD) 0.034 T - 2,400,0000.068 (JD) 0.032T - 2,400,000 (JD) Figure150 1. The combined orbits of the three SB2 observed with PIONIER. Upper row: the40 visual orbits; the node line is in dashes. 4.1 Derivation of the masses M1 (M⊙)0.9820.98081.686Second row: the residuals along the semi-major axis150 of the error ellipsoid. Third row: the residuals along the semi-minoraxisoftheerror ellipsoid.100 Last row: the spectroscopic orbits; the circles refer to the primary component, and the20 triangle to the secondary; the large filled 100 symbols refer to the RV measurements± obtained0.010 at OHP with the±Sophie0.0040spectrograph. The RV0 are± all0.021 shifted to the zero point of the ancient50 measurements of the primary component. The masses of the components are directly derived from the 50 -20 M2 (M⊙)0.88190.72691.3900 RV (km/s) -40 the estimator of the goodness-of-fit defined in Stuart0 & Ord used simulations to verify that this property is true even interferometric and from the RV measurements, taken into -50 ± 0.0089 ± 0.0019 ± 0.034 (1994): when the number of-60 degree of freedom is as small as five. account simultaneously. However, we increase the RV un- ϖ (mas)0 0.5 19.557 1 0 16.702 0.5 1 0 8.445 0.5 1 Phase Phase Phase 1/3 Achieve9ν 1/2 χ 2±1%0.078 or2 better precision± 0.037 on stellar± 0.075 mass! certainties of HIP 117186 by 1.088, which would lead to a Figure 1. TheF = combined orbits of the three+ SB2− 1 observed with(1) PIONIER. Upper row: the visual orbits; the node line is in dashes. 2 2 ν 9ν × Second row: the residuals! " along#! the" semi-major axis$ of the error ellipsoid. Third row: the residuals along the semi-minoraxisoftheerror SB2 orbit with F2 =0.Thisoperationincreasestherelative nVLTI 2201214ellipsoid. Last row: the spectroscopic orbits; the circles refer to the primary component, and the triangle to the secondary; the large filled symbolswhere referν tois the the RV number measurements of degrees obtained of freedom at OHP andwithχ the2 Sophie spectrograph.It appears that, The whenRV are the all uncertainties shifted to the provided zero point by of the the σ(o−c) VLTIancient(mas) measurements of the primary 0.035 component. 0.031 0.0084 weights of the interferometric measurements in the deriva- is the weighted sum of the squares of the differences be- PIONIER reduction are taken into account, F2 is system- n tween the predicted and the observed23+8 values, normalized 63+12atically negative. This indicates 19+7 that these uncertaintiesare RV 1 with respect to their uncertainties. When the predicted val- overestimated. In order to keep the relative weights of the tion of the combined orbit. The solution consists in up to 13 the estimator of the goodness-of-fit defined in Stuart & Ord used simulations to verify that this property is true even ues are obtained− through1 a linear model, F follows the nor- observations, the uncertainties of the positions of any star σ(o−c) RV 1(1994):(km s )0.562,0.1560.952,0.0152.40,0.892 when the number of degree of freedom is as small as five. independent parameters, which are: the orbital parameters mal distribution N (0, 1). When non–linear models are used, are divided by the same coefficient, in order to have F2 =0. nRV 2 but when the errors are small in23+8 comparison to the mea- 4+12The uncertainties σmin and 19+7σmax in Table 1 are thus ob- surements, as hereafter, the model is approximately linear tained. It is worth noticing that they are similar to the errors P , T0, e, V0, ω1, i, Ω1,themassesM1, M2,thetrigonomet- 1/2 2 1/3 around the solution,−19ν and F χfollows also2 N (0, 1). We also expected for Gaia (Perryman 2005; Eyer et al. 2015). σ(o−c) RV 2 (kmF s2 = )0.646,0.2822.54,0.1431.11,0.232 + − 1 (1) a n 2 # ν 9ν $ ric parallax ϖ,andalsotheRVoffsets dn−a, d2−1 and d2−1. −1 ! " ! " dn−a (km⃝c s2015 RAS,)0.323-0.4410.822 MNRAS 000,1–10 2 It is worth noticing that we prefer to directly obtain M1, where ν is the number of degrees of freedom and χ It appears that, when the uncertainties provided by the is the weighted sum of the squares± 0.172 of the differences be- ± PIONIER0.111reduction are taken± 0.563 into account, F2 is system- M2 and ϖ,ratherthantheobservationalparametersK1, K2 a tween−1 the predicted and the observed values, normalized atically negative. This indicates that these uncertaintiesare d2−1 (kmwith s respect)0.4080.4360.549 to their uncertainties. When the predicted val- overestimated. In order to keep the relative weights of the ues are obtained through a linear model, F2 follows the nor- observations, the uncertainties of the positions of any star and a,theapparentsemi-majoraxisoftheinterferometric mal distribution N (0, 1). When± non–linear0.198 models are used, ± are1.513 divided by the same coe±fficient,0.606 in order to have F2 =0. n but− when1 the errors are small in comparison to the mea- The uncertainties σmin and σmax in Table 1 are thus ob- orbit. The advantage of this method is that it leads directly d2−1 (kmsurements, s )-0.1490.077-0.184 as hereafter, the model is approximately linear tained. It is worth noticing that they are similar to the errors to the uncertainties of the masses and of the parallax, in around the solution, and F2 follows± 0.161 also N (0, 1). We also ± expected0.135 for Gaia (Perryman± 2005;0.335 Eyer et al. 2015). ⃝c 2015 RAS, MNRAS 000,1–10 place of the uncertainties on K1, K2 and a when the latter parameters are obtained from the combined interferometric and spectroscopic observations. The parameters of the com- a the uncertainty refers to the VB solution bined solutions are presented in Table 5. The uncertainties of masses range between 0.0019 and 0.034 M⊙,andtherel- ative errors range between 0.26 and 2.4 %. This is similar to the accuracies expected using the Gaia astrometry.

4.2 Notes on individual objects the components is only 0.090 mas, corresponding to 0.99 HIP 14157. This system is widely discussed by Fekel, solar radius. Since Fekel et al. estimated the stellar radii Henry & Alston (2004), who pointed out that the primary is R1 =0.99 R⊙ and R2 =0.76 R⊙,thesystemiscertainly aBYDravariablestarwithavariabilityamplitudearound an eclipsing one. 0.02 mag. Due to an inclination almost edge-on, the masses HIP 20601. The star is a candidate member of the of the components are close to the minimum masses that Hyades cluster (Perryman et al. 1998; de Bruijne, Hooger- they found. We confirm then that the mass of the K2-K3 V werf & de Zeeuw 2001). Griffinetal.(1985)estimatedthat secondary component is 0.882 M⊙.Thisislargerthanthe the spectral types of the components are probably G6 and canonical value, which is between 0.67 M⊙ (for a K5 V K5, in good agreement with our estimates of the effective star), and 0.79 M⊙ (for a K0 V star) according to Cox temperatures. As a consequence, the masses are around 7 % (2000). Such discrepancy is not surprising, but known since larger than the canonical values listed in Cox (2000). awhile,sinceGriffinetal.(1985)andreferencesthereinal- HIP 117186. The effective temperatures of the compo- ready pointed out that the real masses of K-type stars are nents correspond to spectral types around F5, and the usually 15 % larger than the canonical values. In a similar canonical masses are around 1.4. This well corresponds to way, we find that the primary component is too heavy for the mass of the secondary component, but is around 20 % aK0Vstar.Theminimumprojectedseparationbetween percent less than the mass of the primary component.

⃝c 2015 RAS, MNRAS 000,1–10 GRAVITY

Four-beam combiner

2.0–2.4 µm (K-band)

3-mas resolution

Spectroscopy R ~ 22, 500, and 4500 Main science: Study of stars in Galactic Centre The wind of Eta Carinae

Br Y He I

complex morphology of the wind

Gravity Collaboration 17 La Silla Paranal instruments

UVES SPHERE SOFI FORS KMOS MUSE

SINFONI VISIR GRAVITY EFOSC2

FLAMES PIONIER

VIRCAM + HAWKI OMEGACAM NACO

ESPRESSO HARPS XSHOOTER Happy Observing!