Young planets embedded in circumstellar disks

Young planets embedded in
circumstellar disks
Sascha P. Quanz (ETH Zurich)
Image credit: ESO/L. Calçada
The National Centres of Competence in Research (NCCR)
are a research instrument of the Swiss National Science Foundation
Where, when and how do
(gas giant) planets form?
Gas giant planets are found over a broad range of separations
From radial velocity surveys
From direct imaging
HR8799
HD95086
Marois et al 2010; Rameau et al. 2013a,b; Kuzuhara et al. 2013
GJ 504
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
Gaps in the HD169142 protoplan
(Sub-mm) interferometry
NIR scattered light imaging
Distance (AU)
−150 −100 −50
0
50
100
150
N
a)
1"
Distance (arcsec)
E
0.5"
0"
Dip
−0.5"
AO feature
−1"
−1"
−0.5"
0"
0.5"
Distance (arcsec)
1"
−1"
−0.5"
1"
0"
0.5"
The Astrophysical Journal Letters,
N 729:L17 (6pp), 2011 March 10
c)
150
E
100
Dec. offset (arcsecond)
Distance (AU)
0
0.5
1
100
50
0
-50 -100 -1
50
0
−50
−100
AO feature
−150
100
150
-0.5
−150 −100 −50
0
50
Distance (AU)
Andrews et al. 2011; Hashimoto et al. 2011; 2012;
Quanz et al. 2013b; Havenhaus et al. 2014; Garufi et al. 2013
Fig. 1.— NACO/PDI observations of HD169142 in the H
stellar flux (image shown in a linear stretch). The position
E of t
regions have been masked out. Our data reveal a bright inner
dip in the ring and a residual AO feature are indicated by arrow
0.5
0
-0.5
-1
image also scaled with r 2 . Features from the AO system and
offse
The innermost, masked out region is less than 0.1′′R.A.
in diamete
10
(mJy/
3 systems with promising
candidate planets in disks
(at the moment)
The planet candidate in the LkCa 15 disk
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
Kraus & Ireland 2012; Andrews et al. 2011; see also Thalmann et al. 2011, 2014, 2015
First attempts to detect the circumplanetary disk around LkCa 15 b
VLA 7mm data
Isella et al. 2014
HD169142 - sequential+planet formation?+
0
DEC offset (arcsec)
1.6 micron
scattered light image
+
−0.4
7 mm
(a)VLA data
(b)
−0.8
0.8
+
VLA CnB+B+A
7 mm
VLA 7 mm
VLT H−Band
0.4
inner
inner gap
gap
(cavity)
(cavity)
+
0
+
ring
ring
29 AU
outer gap
outer gap
−0.4
−0.8
(c)
0.8
•Inner cavity <25 AU
•Annular gap ~40-70 AU
Quanz et al. 2013b; Osorio et al. 2014
??
(d)
0.4
0
•
−0.8 0.8
0.4
−0.4
RA offset (arcsec)
0
−0.4
−0.8
Fig. 1.— VLA images of the 7 ~5
mm dust
thermal
emission in several array configurations.
sigma
‘overdensity’
Panels (a) and (b) show, respectively, the CnB and B configuration images. Panel (c) shows
inside the cavity ~50 AU
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
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
Quanz et al. 2013b; Reggiani, Quanz et al. 2014 (see also, Biller et al. 2014)
so that the CO hotband lines have the same equ
B.
2003
in the 2003 spectrum. Table 3 shows the2006
scaled
v = 1–0 emission in the resulting spectrum. W
the 2003 spectrum to obtain the spectrum of th
excess emission component (Figure 4). The EW
CO emission component and its velocity centro
are shown in Table 3 where they are compared
from all earlier epochs.
D.
2010
2013
As described in Paper II and summarized in T
2003 and 2006, the red side of the P26 line bright
offset of the red side of the line decreased, and
emission component had a velocity centroid of
(compare Figures 4(A) and (B)). Similarly, in 20
brightened further, and the excess emission wa
zero velocity (−1 ± 1 km s−1 ). In this part of
the spatial centroid of the line became offset fu
(Figure 4(C)). In 2006
2013
here,
km sepoch reported
=+6±1
vthe
E
φ=47±10
the P26 line is again
brighter than in 2003 and th
PA=140
N
blueshifted (−6 ±PA=−5
1 kmE ofs−1
).2003
The spatial centroi
2010
of the line is comparable to φ=0
that in 2003 (Table
v =−1±1 km s
PA=−40
E of N to the eas
side
of
the
line
is
now
extended
further
φ=97±7
A.
HD100546 - sequential planet formation again?
1.6 micron
0.5" light image
scattered
OH Emission
2010: Black
2013: Red
C.
High dispersed 4.6 micron spectroscopy
(B)
CO Emission
2003: Black
2006: Red
2010: Blue
2013: Green
E.
-1
p
o
o
o
o
p
-1
o
o
PA=60o E of N
2013
vp=−6±1 km s-1
φ=133±10o
PA=105o E of N
4. DISCUSSION
4.1. Orbital Analysis
In Paper II, we suggested that the variatio
1–0 line emission could be explained by the
spatially concentrated source of CO emission
star within the disk wall. A schematic of this sc
0.5"
in Figure 4(E). We can obtain a rough constrain
the CO excess component given the velocity cen
Fig. 4.— Spectroastrometric signal of the P26 line and schematic of the geometry of the system. In Panels A–D the
Figure 3. Multi-epoch observations of the of the OH (panel
andlineCO
signal (A))
of the P26
is plotted. The
flux of the
P26 lineand
is plotted
below the
spectro-astrometric
in Pa
inexcess
2006,
2010,
2013.
Assuming
a signal
system
theobserved
spectro-astrometric
our excitation model with the excess emission added for the dat
(panel (B)) lines. In panel (A), the average of the OH emissionepoch
lines
in signal isIncalculated
◦Panel E, from
2010, and 2013 (red dot-dashed line).
a schematic
of the2007;
disk and extra
emissionet
source
is 2014)
presented. The
42
(Ardila
et
al.
Pineda
al.
and
by the
blackhave
dashed
line. The inner wall of the disk orange. The location of the source of the emission excess is labele
2010 (black) and 2013 (red) are plotted over one another. Both
lines
been
and the uncertainty in the phase of the
orbit
is represented
by the
red a
triangle.
In 2003 we
assume to
the emission
is hidd
of
2.4
M
,
we
fit
circular
orbit
the
meas
⊙
scaled to a constant equivalent width. The difference between
these
spectradisk.
is The phase of the orbit is calculated
of the
circumstellar
from the Doppler shift of the excess emission assuming
and the orbital radius is 12.5 AU
(just inside
inner rim of
the disk).R
In and
2006, the
excess
emission pulled
the ce
capturing the structure of the disk,
with
thetheorbital
radius
the
orbital
phase
plotted above. While the equivalent width of the lines varied,42
the
shape of the
red side of the line closer to the the center of the PSF. In 2010, the excess emission pulled both sides of the line east
axis. In 2013, In
the panel
excess emission on the blue side of the line pulled the spectro-astrometric signal eastward. 2
lines has not varied to within the signal to noise of our measurement.
2003
as
free
parameters.
The
result
of
a
χ
fit (
w the noise signature, and P , which
(B) we plot the overlapping region of the CO spectra observed over four epochs.
+1.5
R = 12.9−1.3 AU and an orbital phase of φ =
spectra have been scaled so that the equivalent width of the average of the
multiple-scattering effects (see The
also
hotband
lines is constant. While the shape of the hotband lines has not changed
φ = 0◦ corresponds to the NW end of the sem
over the four epochs spanning 2003–2013, the v = 1–0 P26 line has varied. In
positions where no data is available
we adopt a higher inclination (e.g., 50◦ ; compa
2006, the P26 line shows a red excess relative to the 2003 spectrum. In 2010,
Quanz et al. 2011 and Panić et al. 2014), our b
excess shows a minimal Doppler shift (−1 ± 1 km s−1 ) relative to 2003. In
e left in all images. The imagestheare
the P26 line shows a blueshifted excess.
R = 14.0+1.6
−1.3 AU and an orbital phase of φ = 1
scale in each of the P images. 2013,
All
(A color version of this figure is available in the online journal.)
Hence, the velocity centroids at the three m
are consistent with circular motion and locate th
50 au
•Inner cavity <14 AU
•Brightness asymmetry
•Fundamental CO ro-vibrational
lines
•Hot-band lines static
•v=1-0 P26 line varies
Avenhaus, Quanz et al. 2014; Quanz et al. 2011; Brittain et al. 2013,2014
•Spectro-astrometric signal
consistent with orbiting body
at ~10-12 AU
4
HD100546 - sequential planet formation again?
1.6 micron
0.5" light image
scattered
High contrast imaging
50 au
0.5"
•Point source + plus extended emission at ~52 AU
•Inner cavity <14 AU
capturing
theasymmetry
structure of the disk,
•Very red; not detected shortward of 3.8 micron (yet)
•Brightness
w the noise signature, and P , which
•Teff ~ 930 K
multiple-scattering effects (see also
•R = 7 RJupiter
positions where no data is available
e left in all images. The images are
scale in each of the P images. All
Avenhaus, Quanz et al. 2014; Quanz et al. 2011, 2013a, 2015; Currie et al. 2014
What’s next? - Get more data…
ALMA cycle 3 simulations (345 GHz)
HD100546
HD169142
Take home messages
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