SCExAO: a platform for high contrast imaging at the Subaru Telescope

SCExAO: a platform for high contrast imaging at the Subaru Telescope

Frantz Martinache, Laboratoire Lagrange, OCA

Olivier Guyon & Nemanja Jovanovic, Subaru Telescope, NAOJ

interferometry coronagraphy

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SCExAO primary WFS: pyramid-based

- OCAM2k (3.5 kHz, 0.3e- readout)- BMC 2k actuator DM

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SCExAO after AO188

November, 2014

SCExAO IR: the tools for wavefront control

- deformable mirror- science camera- access to the pupil

Jovanovic et al, 2015, submitted

The AO conundrum

Atmosphere

Telescope

Instrument

Camera

fast(milli second)

slow(~ 1 hour)

very slow(~ 1 day)

Wavefront sensor

DM

fast(milli second)

HR 8799

|FT|2

Image

ModulationTransferFunction

simple complex

Φ(1-2) = Φ(1-2)0 + (Φ1-Φ2)Φ(2-3) = Φ(2-3)0 + (Φ2-Φ3)Φ(3-1) = Φ(3-1)0 + (Φ3-Φ1)

Jennison, 1958, MNRAS, 118, 276

To take advantage of self-calibration

the closure-phase:

Kraus & Ireland, 2012, ApJ, 745, 5

High contrast detection near the diffraction limit

To what end?

Φ(2-1) = Φ(2-1)0 + (φ2-φ1)Φ(3-2) = Φ(3-2)0 + (φ3-φ2)Φ(1-3) = Φ(1-3)0 + (φ1-φ3)

... ... ... ... ... ...Φ(k-l) = Φ(k-l)0 + (φk-φl)

... ... ... ... ... ...

Φ Φ0= + φA ×

measuredFourierphase

trueFourierphase

phasetransfermatrix

pupilwavefront

errors

Martinache, 2010, ApJ, 724, 464

Φ = ΦO + A.φ

K ϕ = K ϕo + K A φK ϕ = K ϕo(kernel-phase)

Martinache, 2013, PASP, 125, 422

Martinache, 2010, ApJ, 724, 464

φ = A-1 . (ϕ-ϕo)(eigen-phase)

So what?

Pupil asymmetry required

Martinache, 2013, PASP, 125, 422

Application to SCExAO

Instrument model: - 292 pupil vector- 6xx uv points

Reconstruction:SVD: 150/292 modescontrol: 7 Zernike modes

June 2014, on- & off-sky calibration of the NCP error with only a pupil mask

SCExAODM

monitor

SCExAOAPF-WFS

GUI

Even with a perfect wavefront...diffraction everywhere!

Broadband H image with SCExAO

PIAA optics on SCExAO

First on sky demonstration (Sept. 2011)

Martinache & Guyon, 2011, AO4ELT2

- IWA ~ 1 λ/D- Throughput: ~ 100 %

SCExAO’s baseline coronagraph: PIAA

Carve a dark hole in your image

SCExAO cal source, non coronagraphic imageMartinache et al, 2014, PASP, 126, 565

PIAA: compatible with focal plane WF control

Martinache et al, 2012, PASP, 124, 1288

Requirements

Martinache et al, 2014, PASP,126, 565

1. Fast upstream XAO correction2. Fast and sensitive camera for speckle sensing:- SAPHIRA (IfA, Nov 2014)- MKIDS (UCSB, 2016)

In the mean time:- proof of concept- plan the software ahead- integrate the feedback to the upstream WFS

DM shape50-ms exposure coadded exposure

Blind speckle nulling on sky with partial AO

June 2, 2014 observations of RX Boo

Hitch-hiker experiments


Diffraction limited imaging in the visible640 - 680 nm wavelengthAchieved ~17 mas FWHM despite a 2” seeing in the visibleRelies on a Lucky Fourier algorithm

Figure 5.16: Véga (α Lyrae). A gauche, un zoom sur le cœur de la PSF extraite de l’image issuede l’algorithme ISFS à 1% de taux de sélection pour la série de 104 images prises à λ=680nm. Adroite, pour comparaison, une simulation de la PSF parfaite du télescope Subaru, échantillonnéeà 4.94 mas/pixel. Les deux images sont saturées afin de révéler les anneaux d’Airy. Malgré laforte composante de halo, le premier anneau est clairement visible, les deuxièmes et troisièmesanneaux sont aussi discernables.

146

Figure 5.8: Véga (α Lyrae). Exempled’une image prise lors de la série àλ=656nm. Temps d’exposition de 28.6millisecondes. Le champ total de la ca-méra est de 2.53 secondes d’arc.

5.7) qui varie rapidement du simple au quintuple au cours de la séquence d’acquisition

(104 images à 28.7 millisecondes représentent une durée de 4 minutes 47 secondes).

Lors de cette première nuit, la turbulence est relativement élevée et la correction four-

nie par le système AO188 ne permet pas la formation d’un speckle dominant les autres

par son flux (voir la figure 5.8). Les variations à la fois sur la position du photocentre

de l’image et celle du speckle le plus brillant varie largement durant l’acquisition de cette

séquence d’images. Lors de cette acquisition, la différence entre la position du photocentre

et celle du speckle le plus brillant (voir la figure 5.9) ne révèle pas de direction privilégiée.

D’autres séries d’images (je traite ce cas dans la partie 5.5.1) montreront de telles direc-

tions privilégiées : celles-ci influent sur les résultats obtenus par les algorithmes en terme

de qualité d’images.

Afin d’estimer la qualité de chaque image, je calcule le rapport entre le speckle le plus

brillant de l’image et le flux total. Cet estimateur est proportionnel au rapport de Strehlsi la tâche de seeing reste contenue dans l’image. Ce rapport de flux entre le speckle le

plus brillant et le flux total de l’image (voir la figure 5.10) montre la grande variabilité de

la concentration du flux sur un speckle unique : certaines images ont une bien meilleure

concentration et donc un rapport de Strehl bien plus élevé, favorable aux méthodes de

sélection d’images telles que le Lucky Imaging (le rapport varie du simple au triple).

Ces images arborant un haut rapport sont isolées : leurs voisines immédiates ont un

rapport bien plus faible. (Ceci est un indice, sans être une preuve absolue, que le temps

de cohérence atmosphérique est bien plus court que le temps d’exposition employé.)

Cependant, pour cette série d’images, le rapport entre flux du speckle le plus brillant et

flux total reste toutefois dépendant de l’estimation total du flux. Par exemple, l’épisode de

relativement faible flux entre la 4000ième image et la 5000ième par rapport au reste de la

série dans la figure 5.7 est également un épisode où le rapport est en moyenne plus faible

dans la figure 5.10. La raison est probablement due à une faible fuite de photons hors

137

28 ms exposure of α LyraeNarrow-band filter (10 nm)

@ λ=680 nm

FOV = 2.5”

Diffraction limited image reconstructed by the algorithm

PhD dissertation by Vincent Garrel (Oct 2012)Garrel et al, 2012, PASP, 124, 861

PerformancesParamètres λ=656nm λ=680nm

Sélection Algorithme Rapport Gain FWHM Rapport Gain FWHMde Strehl (%) (mas) de Strehl (%) (mas)

25%Lucky Im. 0.10 2.0 217 0.11 2.3 143ISFAS 0.12 2.3 84 0.13 2.7 79ISFS 0.29 5.4 32 0.35 7.1 20



Table 5.3: Bételgeuse (α Orionis). Caractéristiques des PSF dans les images issues des algo-rithmes en fonction du taux de sélection à partir de deux séries consécutives de 104 images, l’uneà λ=656nm, l’autre à λ=680nm. Le gain est donné par le rapport du pic du cœur de la PSFreconstitué par rapport au pic du halo dans l’image moyenne.

Figure 5.23: Bételgeuse (α Orionis). A gauche, un zoom sur le cœur de la PSF extraite del’image de Bételgeuse issue de l’algorithme ISFS à 1% de taux de sélection pour la série de104 images prise à λ=680nm. A droite, pour comparaison, la PSF de Véga également issue del’algorithme ISFS à 1% de taux de sélection sur une série de 104 images à λ=680nm, au mêmeéchantillonnage. Il semble donc que l’algorithme ISFS parvienne à fournir en partie une imagerésolue de Bételgeuse.

154

Resolved image of Betelgeuse

λ=680 nm

PerformancesParamètres λ=656nm λ=680nm

Sélection Algorithme Rapport Gain FWHM Rapport Gain FWHMde Strehl (%) (mas) de Strehl (%) (mas)




Table 5.3: Bételgeuse (α Orionis). Caractéristiques des PSF dans les images issues des algo-rithmes en fonction du taux de sélection à partir de deux séries consécutives de 104 images, l’uneà λ=656nm, l’autre à λ=680nm. Le gain est donné par le rapport du pic du cœur de la PSFreconstitué par rapport au pic du halo dans l’image moyenne.

Figure 5.23: Bételgeuse (α Orionis). A gauche, un zoom sur le cœur de la PSF extraite del’image de Bételgeuse issue de l’algorithme ISFS à 1% de taux de sélection pour la série de104 images prise à λ=680nm. A droite, pour comparaison, la PSF de Véga également issue del’algorithme ISFS à 1% de taux de sélection sur une série de 104 images à λ=680nm, au mêmeéchantillonnage. Il semble donc que l’algorithme ISFS parvienne à fournir en partie une imagerésolue de Bételgeuse.

154

λ=680 nm

β Delphini

0.24”

λ=656 nm

PhD dissertation by Vincent Garrel (Oct 2012)

FIRSTby the group at LESIA (G. Perrin, E. Huby et al)

VAMPIRES

6 B. Norris et al.

To all

To EMCCD

To EMCCD

MaskWheel

HWP

SCExAO optics

To PyWFSPeriscope M2

Periscope M1

BRT L2

Pupil Plane

FilterWheel

TCLCC

Computer

LCVR Wolls.Prism

EMCCDCamera

BRT L1

Cam L1 Cam L2Dicr. M2

Dicr. M1 OAP2 OAP1DM

Translation stageRotation stage

To coronagraphs,IR cameras

Telescope(M1, M2, M3)

QW

P

IP Cal. SourceSu

per-K

OAP

col

.

Dep

olz

ArduinoSMF

Lin

polz

AO 188

Figure 1. A schematic diagram of VAMPIRES as configured on-sky in July 2013, with all items relevant to the VAMPIRES beamtrain shown. Operation of each subsystem is described in the text. Abbreviations: M - Mirror; L -Lens; OAP - Off Axis Parabola; DM -Deformable Mirror; Dicr.M - Dichroic Mirror; HWP - Half-wave plate; BRT - Beam Reducing Telescope; LCVR - Liquid Crystal VariableRetarder; LCC - LCVR Controller; TC - Temperature Controller; Cam - Camera; QWP - Quarter-Wave Plate; Depolz - Depolariser;OAP col. - OAP Collimator; Lin polz - Linear polariser. In an alternative configuration, the half-wave plate can be replaced with a pairof quarter-wave plates to allow birefringence to be corrected as needed

Position on pupil (m) Position on pupil (m) Position on pupil (m) Position on pupil (m)

Posit

ion

on p

upil

(m)

Baseline length (m)

Base

line

leng

th (m

)

0

0

-5

-5

-10-10

5

5

10

10

Baseline length (m)0

0

-5

-5

-10-10

5

5

10

10


0

-5

-5

-10-10

5

5

10

10


0

-5

-5

-10-10

5

5

10

10

Figure 2. The non-redundant aperture mask designs installed in VAMPIRES (top) and their corresponding Fourier coverage (bottom).Masks with a greater number of holes boast better Fourier coverage at the expense of throughput. Masks are all non-redundant (thevector separation of all hole pairs is unique), with the exception of the annulus mask (right). This is doubly-redundant, sacrificing non-redundancy for high throughout and full Fourier coverage. The gaps in Fourier coverage for this mask are due to the missing portions ofthe annulus needed to screen out the secondary-mirror supports (spiders).

passed through an achromatic linear polariser (Thorlabs

LPVIS100-MP) mounted in a CONEX-AG-PR100P rota-

tion stage to allow linear polarised light at any specified

orientation to be injected. A quarter-wave plate (Thorlabs

AQWP10M-980), also in a rotation mount, is also moved

into the beam allowing a complete Mueller matrix of the

instrument to be constructed as described in Appendix A.

As described in Section 3, VAMPIRES employs three

tiers of differential calibration, one of which is rapid channel

switching using the LCVR. This requires alternate frames

acquired by the camera to have the incident beam polarisa-

tion rotated by 90◦, thus swapping the state probed by the

twoWollaston channels. Due to the high acquisition rate and

short integration times used (∼17 ms in order to maintain

high visibilities despite residual seeing after AO correction)

it was not possible to directly control the LCVR switch-

ing and camera exposures using the computer, due to the

non-realtime operating system and variable USB latency.

Instead, the timing signals were generated by a dedicated

Arduino Uno microcontroller. When a data acquisition cycle

c� 2014 RAS, MNRAS 000, 1–14

NRM interferometry + polarimetry = differential measurements

by the group at USyd (P. Tuthill, B. Norris et al)

VAMPIRESThe VAMPIRES instrument 9

! = 6.8 %

a) 1 tier - Wollaston prism only. No temporal variation leads to small error bars, but strong systematic errors (from non-common path) dominate.

! = 2.4 %

b) 1 tier - LCVR only. No non-common path error, and the mean is ~ 1.0. However since switching is slower than seeing temporal errors lead to large error bars.

! = 0.86 %

c) 2 tiers - Wollaston + LCVR. The Wollaston and LCVR cancel each others errors. Systematic errors are still visible.

! = 0.42 %

d) 3 tiers - Wollaston + LCVR + HWP. The HWP cancels out static systematic errors (such as those arising from instrumental effects). Here precision is limited by random error; additional integration time would improve precision further.

0 7.22.4 4.8Baseline Length (m)

Figure 4. The on-sky differential visibilities from an observation

of Vega at 775 nm with the 18 hole mask, showing the effectof different tiers of calibration. Ideally the visibility ratio should

be unity on all baselines, since the source is unresolved. Baseline

azimuth is plotted on the horizontal axis, while baseline length

is represented by colour. The precision is seen to increase with

successive layers of calibration, as discussed in the text. Data

were taken without Extreme-AO correction.

target). However non-polarimetric measurements were made

using the previously observed star Altair (discussed above)

as a PSF reference, although this was not an ideal calibrator

since it was observed at a different air-mass and time of the

night. Despite this, accurate complex visibility data were

recovered, constraining a uniform disk fit yielding a diame-

! = 0.60 %

a) Annulus mask, all baselines. Errors from long baselines due to pupil misalignment dominate.

! = 0.17 %

b) Only baselines shorter than 4 m (unaffected by pupil mislagniment).


Figure 5. The on-sky triple differential visibilities from Vega at

775 nm, with the annulus mask. Due to a misalignment between

the mask and pupil, many longer baselines have extremely low

visibilities, resulting in large errors (panel a). If these affectedbaselines are eliminated by only plotting shorter baselines, excel-

lent precision (0.17 %) is observed (panel b).

ter of 32.2 ± 0.1 milliarcseconds. This is in close agreement

with the literature values tabulated in the CHARM2 catalog

(Richichi et al. 2005), which gives the uniform disk diameter

as 32.8 ± 4.1 milliarcseconds in V band.

The binary system η Pegasi was observed with the 18

hole mask at λ = 775 nm, again for a total integration time

of 54 s. Vega was again used as a calibrator (with the same

reservations). The binary was detected, and its separation

and position angle constrained. A Monte Carlo simulation

was used to determine the statistical confidence of the de-

tection, which was found to be better than 99.9%. The sepa-

ration was measured to be 48.9 ± 0.6mas. This is consistent

with the predicted separation based on the orbital param-

eters measured by Hummel et al. (1998) of 49.9mas. The

slight discrepancy is probably a result of imperfect knowl-

edge of the mapping between the sky and instrumental field

orientations, which is presently based only on values from

the optical system model. Further studies of several stellar

systems with known structure are planned to precisely cal-

ibrate both orientation and plate scale of VAMPIRES. The

contrast ratio was measured to be 3.55 ± 0.06 magnitudes,

again in good agreement with the value measured by Hum-

mel et al. (1998) of 3.61 ± 0.05 magnitudes.

5 SIMULATED DATA AND PERFORMANCE

PREDICTIONS

The differential Fourier visibilities (e.g. VHoriz/VVert) ob-

tained from VAMPIRES are not directly equivalent to the

differential intensities (or fractional polarisations) obtained

in techniques such as polarised differential imaging. Rather,

c� 2014 RAS, MNRAS 000, 1–14

The VAMPIRES instrument 9

! = 6.8 %

a) 1 tier - Wollaston prism only. No temporal variation leads to small error bars, but strong systematic errors (from non-common path) dominate.

! = 2.4 %

b) 1 tier - LCVR only. No non-common path error, and the mean is ~ 1.0. However since switching is slower than seeing temporal errors lead to large error bars.

! = 0.86 %

c) 2 tiers - Wollaston + LCVR. The Wollaston and LCVR cancel each others errors. Systematic errors are still visible.

! = 0.42 %

d) 3 tiers - Wollaston + LCVR + HWP. The HWP cancels out static systematic errors (such as those arising from instrumental effects). Here precision is limited by random error; additional integration time would improve precision further.


Figure 4. The on-sky differential visibilities from an observation

of Vega at 775 nm with the 18 hole mask, showing the effectof different tiers of calibration. Ideally the visibility ratio should

be unity on all baselines, since the source is unresolved. Baseline

azimuth is plotted on the horizontal axis, while baseline length

is represented by colour. The precision is seen to increase with

successive layers of calibration, as discussed in the text. Data

were taken without Extreme-AO correction.

target). However non-polarimetric measurements were made

using the previously observed star Altair (discussed above)

as a PSF reference, although this was not an ideal calibrator

since it was observed at a different air-mass and time of the

night. Despite this, accurate complex visibility data were

recovered, constraining a uniform disk fit yielding a diame-

! = 0.60 %

a) Annulus mask, all baselines. Errors from long baselines due to pupil misalignment dominate.

! = 0.17 %

b) Only baselines shorter than 4 m (unaffected by pupil mislagniment).


Figure 5. The on-sky triple differential visibilities from Vega at

775 nm, with the annulus mask. Due to a misalignment between

the mask and pupil, many longer baselines have extremely low

visibilities, resulting in large errors (panel a). If these affectedbaselines are eliminated by only plotting shorter baselines, excel-

lent precision (0.17 %) is observed (panel b).

ter of 32.2 ± 0.1 milliarcseconds. This is in close agreement

with the literature values tabulated in the CHARM2 catalog

(Richichi et al. 2005), which gives the uniform disk diameter

as 32.8 ± 4.1 milliarcseconds in V band.

The binary system η Pegasi was observed with the 18

hole mask at λ = 775 nm, again for a total integration time

of 54 s. Vega was again used as a calibrator (with the same

reservations). The binary was detected, and its separation

and position angle constrained. A Monte Carlo simulation

was used to determine the statistical confidence of the de-

tection, which was found to be better than 99.9%. The sepa-

ration was measured to be 48.9 ± 0.6mas. This is consistent

with the predicted separation based on the orbital param-

eters measured by Hummel et al. (1998) of 49.9mas. The

slight discrepancy is probably a result of imperfect knowl-

edge of the mapping between the sky and instrumental field

orientations, which is presently based only on values from

the optical system model. Further studies of several stellar

systems with known structure are planned to precisely cal-

ibrate both orientation and plate scale of VAMPIRES. The

contrast ratio was measured to be 3.55 ± 0.06 magnitudes,

again in good agreement with the value measured by Hum-

mel et al. (1998) of 3.61 ± 0.05 magnitudes.

5 SIMULATED DATA AND PERFORMANCE

PREDICTIONS

The differential Fourier visibilities (e.g. VHoriz/VVert) ob-

tained from VAMPIRES are not directly equivalent to the

differential intensities (or fractional polarisations) obtained

in techniques such as polarised differential imaging. Rather,

c� 2014 RAS, MNRAS 000, 1–14

multi-step calibration of the visibilities

Norris et al, 2015, MNRAS, 447, 2894

The VAMPIRES instrument 11

!"

#" $"


Figure 7. A modelled protoplanetary disk at 500 pc (see Section 5) and the derived VAMPIRES data for λ = 800 nm. a) Image of

the inner region of the disk, shown with a non-linear intensity mapping, in four polarisations: horizontal H and vertical V corresponding

to Stokes Q, while H45 and V45 are the two orthogonal polarisations rotated 45◦corresponding to Stokes U. b) The differential power

spectra for the two pairs of orthogonal polarisations. c) Expected differential visibility signals as seen by VAMPIRES with an 18 hole

mask.

between 25% and 100% of the integration time dependingon switching frequency.

The representative model presented in this section hasdifferential visibilities with an average magnitude (devi-ation for unity) of approximately 2%. With the demon-strated on-sky precision using the 18 hole mask of 0.4%,the VHoriz/VVert of each baseline can be measured to 5σ.However the actual uncertainties on a fitted model wouldbe much smaller due to the relatively small number of freeparameters involved.

6 SUMMARY

By combining non-redundant aperture-masking interferom-etry with differential polarimetry, the VAMPIRES instru-ment will directly image the inner-most region of protoplan-etary disks, providing critical insight into the processes ofdisk evolution and planet formation. Non-redundant aper-ture masking provides diffraction-limited performance byway of the established interferometric visibility and closure-phase observables. VAMPIRES’ triple-differential polari-

metric calibration strategy exploits the polarisation of scat-tered starlight, utilising simultaneous differential measure-ments with a Wollaston prism, fast channel-switching witha liquid-crystal variable retarder and slow-switching witha rotating half-wave plate to better remove instrumentalsystematics. The resulting signal encodes the resolved, po-larised structure of the inner disk. These observables arelargely immune to the effects of instrumental polarisation,with the remainder being removed by precise calibration ofthe instrumental Mueller matrix using an in-built character-isation system.

VAMPIRES records data at visible wavelengths (wherepolarisation from scattering is typically higher) in a hitch-hiker mode that is invisible to simultaneous science oper-ation of other (infrared) instruments. On-sky demonstra-tions of the VAMPIRES instrument yielded a differential-visibility precision approaching 10−3 and closure phasestandard-deviation better than 1◦. Limitations to both per-formance metrics are presently provided by restricted sta-tistical sample size and therefore further improvement is ex-pected with longer on-sky integration times. Precise visibil-ities and closure phases will be used to accurately constrain

c� 2014 RAS, MNRAS 000, 1–14

VAMPIRES science case

Norris et al, 2015, MNRAS, 447, 2894

Interferometry & coronagraphy working together: a solution?

Simultaneous WF control and high contrast imaging

Combine the benefits of contrast (coronagraphy) and calibration (interferometry) for high contrast

DM or diffractive grid

Lyot stop+ asymmetry

coronagraph

WFS on satellitespeckles

SCExAO: a platform for high contrast imaging at the Subaru Telescope

Frantz Martinache, Laboratoire Lagrange, OCA

Olivier Guyon & Nemanja Jovanovic, Subaru Telescope, NAOJ


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SCExAO: a platform for high contrast imaging at the Subaru Telescope

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