Near Infrared Spectrometers and

Integral field spectroscopy

Andy Longmore (UK ATC)

(contact via supa@roe.ac.uk)

• Recap of principles of IR instrument design• Some spectroscopy related specifics• Introducing the dispersion element• Spectrometer design - slit width, resolution and instrument size• Novel dispersers• Integral field and multi-object spectroscopy

Outline of lecture

Principles of IR instrument design• As for imaging, control of the infrared background is key

• cryogenic temperatures for NIR (130K-77K) •use of cold pupil stop• conservation of etendue/throughput from telescope to detector

• Challenge remains to detect a faint source against a bright background

• 5sigma in 1 hour for magnitude 17 – 18.5• broadband backgrounds 13-16.5

• Nature of the background very different…..

NIR spectroscopy basics

The OH emission spectrum:

OH spectrum from Rousselot et al. (2000, A&A, 354,1134).

• Lines are bright– 2500 photons per

second per arcsec with an 8-m telescope, 1000 times target source flux

• Variable• BUT, the continuum

between them is very dark– 0.1photons per second

Resolving out the OH background

From Martini & DePoy 2000, Proc SPIE 4008, 695

Spectral resolving powers of >2000, preferably around4000 are a good choice.

Slit alignedon object

Collimator DispersingElement

Camera

X

Y

Y

Detectorarray

entrance aperture/slit + collimator + dispersing component (if possible at a pupil)

Basic Spectrometer layout

Introducing the dispersion element

Grating equation Angular dispersion

Spectral resolving power, R

Slit alignedon object

Collimator DispersingElement

Camera

X

Y

Y

Detectorarray

Angular separation of two wavelengths at the detector

From conservation of étendue:

Wavelengths are resolved when these are equal, so:

NB: max R depends on Pupil size, telescope D and angular size of the object.

Spectral resolving power, R =

Camera focal length, f

Pixel size, w

fw /)tan(

Relating R to instrument design parameters:

)tan(2

rw

fR

NB: often w = two pixels for Nyquist sampling of the spectral line width

KMOS - A practical example

The K-band multi-field spectrograph for the European Southern Observatory Very Large Telescopes

KMOS - A practical example

• Conservation of A and the detector pixel size determines the camera f-number :

• FOV = field of view per pixel in arcsec

f_number

hpixel_widt

206265

FOVD

detectortel

KMOS - a practical example

KMOS - Design choices: pixel FoV

• pixel FoV depends on– Size of images delivered to the

telescope

– the science case

• For maximum sensitivity – Match slit to seeing

• KMOS has designed

for the best seeing at VLT• Pixel scale selected is

0.2arcsec for Nyquist sampling of the seeing

KMOS - Camera f-ratio

• For VLT, D = 8m, and for KMOS pixel width = 18m which we want to match to FOV=0.2arcsec.

• The camera required is f/2.32• NB: a fast camera and difficult to design• Combination of large telescopes and sharp

image through AO provides design changes in future e.g. 40m telescope would require f/0.5 for 0.2arcsec pixels!

• Key requirements for KMOS on spectral resolving power:– high enough to resolve OH lines – Ideally, coverage of a complete transmission

window simultaneously– Allowing velocity resolution of 10km/s for

experiments on the Tully-Fischer relation in galaxies

KMOS - Design choices: gratings

KMOS - Resolving the OH background

From Martini & DePoy 2000, Proc SPIE 4008, 695

• Key requirements for KMOS on spectral resolving power:– high enough to resolve OH lines – Ideally, coverage of a complete transmission

window simultaneously– Allowing velocity resolution of 10km/s for

experiments on the Tully-Fischer relation in galaxies

KMOS - Design choices: gratings

• Atmospheric windows

KMOS - Design choices: gratings

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wavelength (um)

tran

smis

sio

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• Key requirements for KMOS on spectral resolving power:– high enough to resolve OH lines – Ideally, coverage of a complete transmission

window simultaneously– Allowing velocity resolution of 10km/s for

experiments on the Tully-Fischer relation in galaxies

KMOS - Design choices: gratings

• Key requirements for KMOS on spectral resolving power:

– high enough to resolve OH lines

– Allowing velocity resolution of 10km/s for experiments on the Tully-Fischer relation in galaxies

KMOS - galaxy dynamics

• Key requirements for KMOS on spectral resolving power:– high enough to resolve OH lines – Ideally, coverage of a complete transmission

window simultaneously– Allowing velocity resolution of 10km/s for

experiments on the Tully-Fischer relation in galaxies

– Resulting choice: / > 3000, maximum 4000 at J band.

KMOS - Design choices: gratings

• For a grating spectrometer :

– f is the focal length of the camera, – w is the size of the slit at the detector (two 18m

pixels for KMOS) – b the blaze angle for the grating (45º).– f ~ 70mm and beam size ~31mm for R=4000– In practice, KMOS beam size is 33mm

KMOS - gratings again

w

rfR

)tan(2

• Advantages– instrument layout– compact– simplified

Novel Dispersers 1: Grisms

Cam

era Module

Slit Module

Grism

Module

TSIU

Novel Dispersers 1: Grisms

Refraction in the prism:

Diffraction from the grating:

Resulting ‘grism’ equation:

NB: dependence on the refractive index of the glass

Novel Dispersers 1: Grisms

• Advantages– Compact– simplified instrument layout– Potential for higher R – for a given beam size

• Disadvantages: – difficult materials– Relatively low transmission

ND2: Volume Phase Holographic gratings

IR Laboratory prototypes

Lab testing and instrumentdevelopment underway for these gratings.

Real spectrometer throughput

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atmosphere

optimal

essential

Data courtesy UK Infrared Telescope, Tom Ray and Chris Davis

• With typical plate scales of 0.2arsec/pixel and 2048 x 2048 pixel detectors, modern spectrometers have slit lengths of 120arcmin.

The long slit: strengths/weaknesses

An image of a field of stars and galaxies is used to set positions for slitlets in a mask…

the multi-object slit mask is inserted at the telescope focal plane…..

and spectra of the selected sources measured at the detector.

Images from the Gemini Multi-object Spectrograph

The long slit versus MOS spectroscopy

•The distribution of 200,000 galaxies in the 2dF galaxy redshift survey•2dF is an optical multi-object spectrograph with a two-degree field of view and the ability to observe 400 objects simultaneously

A fibre-fed MOS spectrograph: 2DF

Techniques for 3D spectroscopy

Image: Stephen Todd, ROE and Douglas Pierce-Price, JAC

Integral Field Spectroscopy

stray light control

scattering from diamond-machined surfaces

scattering and diffractionfrom edges of the slicing mirror

triple-fold mirror

field andpupil baffle

input baffle

pupil baffle

Dekker mask

mask over unusedmirror surface

IFU

folded optical layout

compact - to fit on slit wheel

input focal plane and outputslit plane are equivalent

f / 10 input and output

deployable

IFU

externalappearance

reality

partiallydismantled

triple-foldmirror

f/conversionmirrorassembly

slicing mirrorassembly

main body

• 0.95mm thick Aluminium slices

• Individually manufactured• Diamond machined• Common radius of

curvature• Surface roughness

~10nm rms• ~700 spatial elements

Metal slicing mirrors

Integral field spectroscopy: data

IFU data

Dispersion

Spa

tial p

ositi

on

yx

Emission line images in 0.4 arcsec seeing

N

E

350 pc

unresolvedSpatially resolved

Wilman, Edge and Johnstone (2005)

IFU data: NGC1275

H2 located in a disk. Resolved for the first time.

ΔV=345 km/s with r=50 pc → MBH = 3.4 x 108 M

A BLACK HOLE MASS ESTIMATE

Techniques for 3D spectroscopy

IFUs 1: SAURON pupil slicing

• SAURON optical IFU on the William Herschel Telescope

• 33’’ x 41’’ total field of view with 0.94” per pixel• Fixed spectral format 4810-5400A (science case…..)

The pupillenslet array

• CIRPASS – Cambridge University

IFUs 2: Fibre bundles for image slicing

Now and in the future: Multi-IFUs

In the future: Multi-IFU systems

MOMSI for a 42-m E-ELT

Precision Radial Velocity Spectrometer: Optical Layout

• Input slit– Output from fibre slicer– 0.9 x 0.08mm effective size, f/5

• Focal reducer– 2x spherical lenses– Convert from f/5 to f/13

• Single pupil mirror– Parabola, f=2000mm, 500x446mm– Used in triple pass

• 140mm collimated beam diameter

• Spectrum mirror– Flat, 386 x 30mm

– Echelle• 31.6 lines/mm, R4 (75° blaze angle)• Small gamma angle (0.4°)• 610 x 200mm

– Cross disperser• Reflective grating, 100 lines/mm, m=1• 210mm diameter

– Camera• Refractive, no aspherics, 6 elements• f=437mm, f/2.78• 0.148 arcsec/pixel

– Detector• 2 x 2K2 HAWAII-2RG arrays

Spectral Format

Detector array footprint2 x 2K2 HAWAII-2RG arrays73.728 x 36.864mm

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Detector X Position (mm)

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ecto

r Y

Po

siti

on

(m

m)

Order 35

Order 611.0m

1.75m

Realistic PRVS Simulations

M6VTeff = 2800 KLog g = 5v sin i = 0 km/s

ModelTelluricOH

Wavelength calibration

(Emission line density used in baseline simulations)

Combined Ar, Kr, Ne and Xe lamps

With Dyanamic range 100-300

Functions of Calibration Unit

• Track displacements at the sub pixel level and changes in Spectral Response Function (SRF)

• Allow accurate wavelength calibration for each order

• Provide Flat Field signal to the spectrograph using a Quartz-Halogen lamp

• Provide a gas cell reference to provide for tracking emission line lamps– Final calibration accuracy must be maintained over

timescales from hours to 5 years

Need large # of lines throughout spectrum

Combine Multiple

emission line lamps in fibers

And: gas absorption cell for additional long-term referenceOR: Laser frequency comb (new technology)

M6V: v vs v sini

John Rayner (University of Hawaii) + PRVS team

Standard Spectroscopy Data Reduction

• Preparation of Single Frames – Manipulation of Raw Data – Preliminaries – Bad pixels

• Creating a bad-pixel mask – Readnoise Variance

• CGS4 Readnoise • Michelle Readnoise • UIST Readnoise • IRIS2 Readnoise • ISAAC Readnoise

– Bias Subtraction – Poisson Variance – Chopping – Flat Fielding – Interleave and Coadd – Orient Image Normally – Wavelength Calibrate

• Group Formation – Sky Subtraction – Group Coaddition

• Spectrum Extraction – Counting Beams – Finding Beams – Beam Extraction – Derippling – Beam Cross Correlation – Extracted Beam Coaddition

• Signal-To-Noise Calculation • Division By Standard Star • Flux Calibration • Science Target Recipe Details • POINT_SOURCE • EXTENDED_SOURCE

– EXTENDED_SOURCE_WITH_SEPARATE_SKY

• FAINT_POINT_SOURCE • POINT_SOURCE_POL • Calibration Recipe Details

http://www.oracdr.org/docs/sun236.htx/sun236.html

3D spectroscopy data processing

3D spectroscopy DR – more examples

• Reduction steps– Wavelength calibrations– Spectral flat-fielding– Fringing– Spatial flat-fielding– Spatial calibration– Sky subtraction– Telluric absorption– Flux calibrations– Cosmic rays

Wavelength Calibration

• What wavelengths fall on which pixels?• How can we re-sample the data to have

linear wavelength axis?– Find dispersion function: relationship between

your pixels and absolute wavelength

• Re-sampling– A necessary evil, at least in spectral direction– Propagates ‘artefacts’ from a sinlge spectrum

to all final spectrum, but diluted so harder to identify.

Spatial Flat-Fielding

• Like an imager, IFUs need 2D spatial flat field• Optical train is much more complex than an

imager.• Vignetting typically due to optics, pupil masks,

etc.• Typically spatial resolution is too low to be very

sensitive to (e.g.) dust particles• Still need global illumination function• Best to use integrated light from the sky rather

than internal lamp

Spatial Flat-Fielding

Lamp

OASIS illumination

Sky

Spatial Calibration

• IFUs have complex optics that may introduce non-negligibledistortions in the spatial plan.

• Important for astrometric applications - in practice, compromisebetween FoV and spatial sampling reduces the significance

• Lenslet-based instruments have spatial stage fixed by lens array -usually well known

• Slicer-based instruments more prone to flexure, esp. along the slit -needs calibrated

• Very important for large-format slicers, like MUSE (ESO)

Reference image(NIRI, Gemini North) NIFS Reconstructed

Sky subtraction

Atmospheric dispersion => ADC(not so important for infrared unless working at diffraction limit)

Atmospheric refraction => image shifts as function of wavelength• Shifts largest at blue wavelengths• Can be corrected during reduction by shifting back each image plane

Co-Adding Data Cubes

• Two approaches:• 1. Dithering by non-integer number of spaxels:

– Allows over-sampling, via ‘drizzling’– Resampling introduces correlated noise– Good for fairly bright sources

• 2. Dither by integer number of spaxels– Allows direct ‘shift and add’ approach– No resampling:- better error characterisation– Assumes accurate (sub-pixel) offsetting– Suitable for ‘deep-field’ applications

3D spectroscopy workshops

• 3D Spectroscopy in Astronomy• Series: Canary Islands Winter School of

Astrophysics• Edited by Evencio Mediavilla • Science perspectives for 3D spectroscopy:

proceedings of the ESO ... Markus Kissler-Patig, Jeremy R. Walsh, … 2007

EAGLE at the

E-ELT Gravity-Invariant Focal Station

Assembly Sequence

The future is HUGE…

Instrument De-Rotator and Services Wrap

Instrument Core

space envelope

Electronics Racks

these can be placedat any point in theassembly sequence

Focal Plate

focal plate is 1226 mmbelow the top of theinstrument envelope

on-axis focus is 1200 mmbelow the top of theinstrument envelope

Receiving Mirrors

mounting plate notyet detailed

TRAMSs and ISSs

ISSs can be attached anddetached individually,following warm-up andback-filling of the cryostat

TRAMS will be mountedwithin an enclosure heldat constant temperature

Robotic Positioner Mounting Ring

Robotic Positioners

Entrance Window

for 5 arcmin science field

LGS Extraction

six LGSs in a ring at7.5 arcmin diameter

Overall Envelope

5m diameter5m high

SideView