The National Centres of Competence in Research (NCCR) are a research instrument of the Swiss National Science Foundation
Young planets embedded in circumstellar disksSascha P. Quanz (ETH Zurich)
Image credit: ESO/L. Calçada
Where, when and how do (gas giant) planets form?
Gas giant planets are found over a broad range of separations
From radial velocity surveys
HR8799
HD95086 GJ 504
From direct imaging
Marois et al 2010; Rameau et al. 2013a,b; Kuzuhara et al. 2013
The physical processes involved in planet formation are largely unconstrained
Spiegel & Burrows 2012 (see also, e.g., Marley et al. 2007)
Indirect signatures of planets thanks to high-spatial resolution imaging of disks
(Sub-mm) interferometry
Andrews et al. 2011; Hashimoto et al. 2011; 2012;Quanz et al. 2013b; Havenhaus et al. 2014; Garufi et al. 2013
The Astrophysical Journal Letters, 729:L17 (6pp), 2011 March 10 Hashimoto et al.
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Figure 2. Magnified view of the inner PI images of AB Aur and their averaged azimuthal profiles. Top: magnified PI image with a coronagraphic occulting mask of0.′′3 diameter (left) and the features of the PI image (right). Central position (0, 0) is the stellar position. The outer and inner rings are denoted by the dashed ellipsoids.The solid ellipsoid indicates the wide ring gap. The dashed circles (A to G) represent small dips in the two rings. The filled diamond, circle, and square representthe geometric center of the inner ring, ring gap, and outer ring, respectively. The field of view in both images is 2.′′0 × 2.′′0. The solid circle in the left bottom insetrepresents the spatial resolution of 0.′′06. Bottom left: averaged azimuthal profiles of the outer ring for the PI (black) and reference PSF-subtracted I (red) images. Theprofile is averaged every 5◦ in position angle (corresponding to resolution) in the outer ring. Bottom right: same with the bottom left image, but for the inner ring withevery 15◦ in position angle (corresponding to resolution) in the inner ring.(A color version of this figure is available in the online journal.)
Our observations are consistent with Perrin et al. (2009).When we assume that a companion is 100% polarized inthe PI image, which is the faintest case as Oppenheimeret al. (2008), the upper limits of its mass at 5σ (the absolutemagnitude of 11.7 at the H band) of the photon noise in DipA are 5 and 6 MJ for an age of 3 and 5 Myr, respectively(Baraffe et al. 2003). These derived upper limits of themasses are consistent with that of 1 MJ inferred by thenumerical simulations (Jang-Condell & Kuchner 2010). Onthe other hand, our upper limits for point sources in the dipsseen in the inner ring are 7 and 9 MJ for these ages due tohigher photon noise.
The structures of AB Aur’s inner (22–120 AU) disk surfacedescribed above indicate that the disk is in an active and probably
early phase of global evolution, and possibly one or more unseenplanets are being formed in the disk.
One possible explanation for the non-axisymmetric structuresis GI of the disk (e.g., Durisen et al. 2007). If Toomre’sQ-parameter (defined as Q = csκ/πG!, where cs, κ , and !are the sound speed, epicycle frequency, and surface density,respectively) is of the order of unity, GI occurs and a modewith a small number of arms is excited, that is, a pattern of thesurface density arises that may resemble what we have observed.However, this GI possibility may be rejected for AB Aur(at present) because optically thin submillimeter observationsindicate that Toomre’s Q-parameter is of the order of 10 (Pietuet al. 2005). It may be noted that the disk mass estimate fromsubmillimeter emission has large uncertainties arising from theuncertainties in the optical properties of the dust particles.
4
NIR scattered light imaging
Gaps in the HD169142 protoplanetary disks revealed by polarimetric imaging 7
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Fig. 1.— NACO/PDI observations of HD169142 in the H band. a) Final Qr image scaled with r2 to compensate for the decrease instellar flux (image shown in a linear stretch). The position of the central star is indicated by the red cross. Saturated pixels in the centralregions have been masked out. Our data reveal a bright inner ring, a large gap and a smooth outer disk in polarized light. A brightnessdip in the ring and a residual AO feature are indicated by arrows. b) Ur with the same scaling and stretch as for the Qr image. c) Intensityimage also scaled with r2. Features from the AO system and the telescope spiders are clearly seen. d) Polar coordinate mapping of Qr.The innermost, masked out region is less than 0.1′′ in diameter. The red line traces the peak brightness of the inner ring.
3 systems with promising candidate planets in disks
(at the moment)
The planet candidate in the LkCa 15 disk
Kraus & Ireland 2012; Andrews et al. 2011; see also Thalmann et al. 2011, 2014, 2015
SMA 850 micron + Keck aperture masking (2.3 and 3.8 micron)
•Dust cavity R~40-50 AU (also in scattered light)•Companion candidate in the cavity at ~11 AU
First attempts to detect the circumplanetary disk around LkCa 15 b
Isella et al. 2014
VLA 7mm data
Quanz et al. 2013b; Osorio et al. 2014
HD169142 - sequential planet formation?1.6 micron scattered light image
•Inner cavity <25 AU•Annular gap ~40-70 AU
– 11 –
inner gap (cavity)
ring
outer gap
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Fig. 1.— VLA images of the 7 mm dust thermal emission in several array configurations.
Panels (a) and (b) show, respectively, the CnB and B configuration images. Panel (c) shows
the image obtained by combining the CnB, B, and A configuration visibilities with a uvrange
<1500 k� (rms=18 µJy beam�1; beam=0.23⇤⇤ ⇥ 0.16⇤⇤, PA=5⇥). Panel (d) shows an overlay
of the image shown in panel (c) (contours) and the VLT/NACO H-band (1.6 µm) polarized
light image from Quanz et al. (2013) (color-scale). Saturated pixels in the central region of
the H-band image have been masked out. In all panels, contour levels are �3, 3, 5, 7, 9, and
11 times the rms. Synthesized beams are plotted in the lower-right corners. The apparent
decrease of the 7 mm emission in the north and south edges of the source is most probably
a consequence of the elongated beam. The larger cross marks the position of the HD 169142
star and the smaller one that of the protoplanet candidate.
7 mm VLA data
• ~5 sigma ‘overdensity’ inside the cavity ~50 AU
Quanz et al. 2013b; Reggiani, Quanz et al. 2014 (see also, Biller et al. 2014)
HD169142 - sequential planet formation?1.6 micron scattered light image
•Inner cavity <25 AU•Annular gap ~40-70 AU
3.8 micron high contrast image
•3.8 micron point source at ~20-23 AU•Not (yet) detected at shorter wavelengths•7mm source not detected
Avenhaus, Quanz et al. 2014; Quanz et al. 2011; Brittain et al. 2013,2014
4 Avenhaus et al.
P� P⇥ P⇥ · 5 P
Hband
Ksband
Hband
Ksband
2006
2006
2013
2013
0.5"
0.5"
0.5"
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Fig. 1.— NACO PDI results in H and Ks filter from epochs 2006 and 2013. From left to right: P⇥, capturing the structure of the disk,P⇤, which is expected to be zero and dominated by noise, P⇤ scaled by a factor of five to better show the noise signature, and P , whichis identical to P⇥ in the absence of any noise and when there is no rotation of the polarization due to multiple-scattering e�ects (see alsotext). Positive values are in orange, negative values in blue. The grey area in the center represents positions where no data is availabledue to saturation e�ects. The red cross marks the position of the star. North is up and east is to the left in all images. The images are1.62�� (� 160 AU) on each side, they all show the same section of the disk. For reference, there is a scale in each of the P images. Allimages scaled with r2.
3.1. Global Scattering Signature
With our new data, we confirm the basic disk structurealready described in (Quanz et al. 2011): The major axisof the disk runs in southeast-northwest direction, andthe brightest parts of the disk are along this axis. Thenortheastern part of the disk appears brighter comparedto the southwestern part. For the first time, we identifya dark lane on this forward-scattering side in all H andKs filter observations including the cube mode observa-tions between �0.2⇥⇥ and �0.6⇥⇥, while the scattered lightpicks up (in this representation scaled with r2) outside
of �0.6⇥⇥.The grains in the disk are preferentially backscattering
in polarization (scattering albedo multiplied with polar-ization fraction, which is what our data measure), whichmakes the far (northeastern) side of the disk appear sig-nificantly brighter. Furthermore, the polarization e⇥-ciency in scattering usually peaks around 90� (e.g., Per-rin et al. 2009), which explains the two bright lobes inthe southeast and northwest: The semi-major axis of thedisk runs along this direction, and the scattering angleat these positions is close to 90� depending on the exact
HD100546 - sequential planet formation again?1.6 micron scattered light image
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D.
Fig. 4.— Spectroastrometric signal of the P26 line and schematic of the geometry of the system. In Panels A–D the spectro-astrometricsignal of the P26 line is plotted. The excess flux of the P26 line is plotted below the spectro-astrometric signal in Panels B–D. For eachepoch the spectro-astrometric signal is calculated from our excitation model with the excess emission added for the data acquired in 2006,2010, and 2013 (red dot-dashed line). In Panel E, a schematic of the disk and extra emission source is presented. The orbit is representedby the black dashed line. The inner wall of the disk orange. The location of the source of the emission excess is labeled with a black dot,and the uncertainty in the phase of the orbit is represented by the red triangle. In 2003 we assume the emission is hidden by the near sideof the circumstellar disk. The phase of the orbit is calculated from the Doppler shift of the excess emission assuming the disk is inclined42� and the orbital radius is 12.5 AU (just inside the inner rim of the disk). In 2006, the excess emission pulled the center of light of thered side of the line closer to the the center of the PSF. In 2010, the excess emission pulled both sides of the line eastward along the slitaxis. In 2013, the excess emission on the blue side of the line pulled the spectro-astrometric signal eastward.
The Astrophysical Journal, 791:136 (7pp), 2014 August 20 Brittain et al.
Table 3Properties of Excess CO v = 1–0 Emission
Date Scaled Equivalent Excess Equivalent Doppler Shift FWHM Position Angle Orbital Phase Red BlueWidtha Width of Excess of Excess of Excess of Excessb Offset Offset
(10−2 cm−1) (10−2 cm−1) (km s−1) (km s−1) (mas) (mas)
2003 Jan 7 4.50 ± 0.14 · · · · · · · · · −40◦ 0◦ 12.7 ± 3.3 14.0 ± 3.32006 Jan 14 5.69 ± 0.59 1.19 ± 0.61 +6 ± 1 6 −5◦ 47◦ ± 10◦ 1.6 ± 4.3 17.8 ± 4.32010 Dec 23 6.39 ± 0.57 1.89 ± 0.58 −1 ± 1 12 60◦ 97◦ ± 7◦ 5.7 ± 4.8 25.9 ± 4.82013 Mar 18 5.89 ± 0.20 1.12 ± 0.15 −6 ± 1 6 105◦ 133◦ ± 10◦ 10.9 ± 0.9 35.8 ± 0.9
Notes.a The spectra were scaled such that the average hot band line profiles observed in 2006 and 2010 had the same equivalent widths as the average line profile observedin 2003.b The phase is measured counter clockwise from the northwest end of the semimajor axis of the disk.
(A)
CO Emission 2003: Black2006: Red2010: Blue2013: Green
OH Emission2010: Black2013: Red
(B)
Figure 3. Multi-epoch observations of the of the OH (panel (A)) and CO(panel (B)) lines. In panel (A), the average of the OH emission lines observed in2010 (black) and 2013 (red) are plotted over one another. Both lines have beenscaled to a constant equivalent width. The difference between these spectra isplotted above. While the equivalent width of the lines varied, the shape of thelines has not varied to within the signal to noise of our measurement. In panel(B) we plot the overlapping region of the CO spectra observed over four epochs.The spectra have been scaled so that the equivalent width of the average of thehotband lines is constant. While the shape of the hotband lines has not changedover the four epochs spanning 2003–2013, the v = 1–0 P26 line has varied. In2006, the P26 line shows a red excess relative to the 2003 spectrum. In 2010,the excess shows a minimal Doppler shift (−1 ± 1 km s−1) relative to 2003. In2013, the P26 line shows a blueshifted excess.(A color version of this figure is available in the online journal.)
vary relative to previous epochs (Figures 3(B) and 4(A)–(D)).We analyze the CO v = 1–0 emission using the procedure andrationale outlined in Paper II. We first normalize the spectrumso that the CO hotband lines have the same equivalent width asin the 2003 spectrum. Table 3 shows the scaled EW of the COv = 1–0 emission in the resulting spectrum. We then subtractthe 2003 spectrum to obtain the spectrum of the CO v = 1–0excess emission component (Figure 4). The EW of the excessCO emission component and its velocity centroid and FWHMare shown in Table 3 where they are compared with the valuesfrom all earlier epochs.
As described in Paper II and summarized in Table 3, between2003 and 2006, the red side of the P26 line brightened, the spatialoffset of the red side of the line decreased, and the CO excessemission component had a velocity centroid of +6 ± 1 km s−1
(compare Figures 4(A) and (B)). Similarly, in 2010 the P26 linebrightened further, and the excess emission was centered nearzero velocity (−1 ± 1 km s−1). In this part of the line profile,the spatial centroid of the line became offset further to the east(Figure 4(C)). In the 2013 epoch reported here, the blue side ofthe P26 line is again brighter than in 2003 and the excess is nowblueshifted (−6 ± 1 km s−1). The spatial centroid of the red sideof the line is comparable to that in 2003 (Table 3), but the blueside of the line is now extended further to the east (Figure 4(D)).
4. DISCUSSION
4.1. Orbital Analysis
In Paper II, we suggested that the variations in the v =1–0 line emission could be explained by the presence of aspatially concentrated source of CO emission that orbits thestar within the disk wall. A schematic of this scenario is shownin Figure 4(E). We can obtain a rough constraint on the orbit ofthe CO excess component given the velocity centroids observedin 2006, 2010, and 2013. Assuming a system inclination of42◦ (Ardila et al. 2007; Pineda et al. 2014) and a stellar massof 2.4 M⊙, we fit a circular orbit to the measured velocitieswith the orbital radius R and the orbital phase of the excess in2003 as free parameters. The result of a χ2 fit (Figure 5) givesR = 12.9+1.5
−1.3 AU and an orbital phase of φ = 6◦+15◦−20◦ , where
φ = 0◦ corresponds to the NW end of the semimajor axis. Ifwe adopt a higher inclination (e.g., 50◦; compare for exampleQuanz et al. 2011 and Panic et al. 2014), our best fit shifts toR = 14.0+1.6
−1.3 AU and an orbital phase of φ = 15◦+13◦−15◦ .
Hence, the velocity centroids at the three measured epochsare consistent with circular motion and locate the excess source
4
High dispersed 4.6 micron spectroscopy
•Fundamental CO ro-vibrational lines•Hot-band lines static•v=1-0 P26 line varies
•Spectro-astrometric signal consistent with orbiting body at ~10-12 AU
•Inner cavity <14 AU•Brightness asymmetry
Avenhaus, Quanz et al. 2014; Quanz et al. 2011, 2013a, 2015; Currie et al. 2014
4 Avenhaus et al.
P� P⇥ P⇥ · 5 P
Hband
Ksband
Hband
Ksband
2006
2006
2013
2013
0.5"
0.5"
0.5"
0.5"
Fig. 1.— NACO PDI results in H and Ks filter from epochs 2006 and 2013. From left to right: P⇥, capturing the structure of the disk,P⇤, which is expected to be zero and dominated by noise, P⇤ scaled by a factor of five to better show the noise signature, and P , whichis identical to P⇥ in the absence of any noise and when there is no rotation of the polarization due to multiple-scattering e�ects (see alsotext). Positive values are in orange, negative values in blue. The grey area in the center represents positions where no data is availabledue to saturation e�ects. The red cross marks the position of the star. North is up and east is to the left in all images. The images are1.62�� (� 160 AU) on each side, they all show the same section of the disk. For reference, there is a scale in each of the P images. Allimages scaled with r2.
3.1. Global Scattering Signature
With our new data, we confirm the basic disk structurealready described in (Quanz et al. 2011): The major axisof the disk runs in southeast-northwest direction, andthe brightest parts of the disk are along this axis. Thenortheastern part of the disk appears brighter comparedto the southwestern part. For the first time, we identifya dark lane on this forward-scattering side in all H andKs filter observations including the cube mode observa-tions between �0.2⇥⇥ and �0.6⇥⇥, while the scattered lightpicks up (in this representation scaled with r2) outside
of �0.6⇥⇥.The grains in the disk are preferentially backscattering
in polarization (scattering albedo multiplied with polar-ization fraction, which is what our data measure), whichmakes the far (northeastern) side of the disk appear sig-nificantly brighter. Furthermore, the polarization e⇥-ciency in scattering usually peaks around 90� (e.g., Per-rin et al. 2009), which explains the two bright lobes inthe southeast and northwest: The semi-major axis of thedisk runs along this direction, and the scattering angleat these positions is close to 90� depending on the exact
HD100546 - sequential planet formation again?
50 au
•Inner cavity <14 AU•Brightness asymmetry
•Point source + plus extended emission at ~52 AU•Very red; not detected shortward of 3.8 micron (yet)•Teff ~ 930 K•R = 7 RJupiter
1.6 micron scattered light image High contrast imaging
What’s next? - Get more data…ALMA cycle 3 simulations (345 GHz)
HD100546 HD169142
For 3 young stars we have observational evidence that young planets might orbit in their disks
If these are indeed forming planets then they are located between ~10-50 AU, i.e., at rather larger separations
The directly imaged planets have very red IR colors indicating the possible existence of warm circumplanetary material
More objects are expected to be found thanks to ongoing ALMA and high-contrast imaging campaigns
These objects allow us to constrain planet formation models with empirical data
Take home messages
