Visible Broadband Imager (VBI) Instrument Science Requirements

Project Documentation Document SPEC-0054

Revision H

Advanced Technology Solar Telescope 950 N. Cherry Avenue Tucson, AZ 85719 Phone 520-318-8102 [email protected] http://atst.nso.edu Fax 520-318-8500

Visible Broadband Imager (VBI)

Instrument Science Requirements

F. Wöger, H. Uitenbroek, A. Tritschler, D. Elmore

Instrument Group August 12, 2010

Name Signature Date

Prepared By: Friedrich Wöger F. Wöger 02-May-12

Approved By: Bill McBride VBI Engineering Manager

W. McBride 03-May-12

Approved By: Simon Craig Systems Engineer

R. Hubbard 25-May-12

Approved By: Thomas Rimmele Project Scientist

T. Rimmele 25-May-12

Released By: Joseph McMullin Project Manager

J. McMullin 25-May-12

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SPEC-0054, Revision H Page ii

REVISION SUMMARY:

Date Changes made By

09/15/2001 Initial version Berger

10/08/2001 Sections 1–4 Berger

10/09/2001 Minor stylistic changes, removed some technical

details Keller

10/16/2001 Added magnetic element section and PD optics Berger

08/04/2003 New dual camera concept and refined optics

discussion Berger

08/15/2003 Added cadence and interface requirements; Refined

optical requirements Berger

01/18/2005 Changed to VLBI name, minor mods Berger

03/31/2006 Changed to VBI. Simplified optical layout to

multiple beam without zoom Uitenbroek & Tritschler

08/08/2006 Minor updates; Release of Revision F. Uitenbroek & Tritschler

04/16/2009 New instrument design, wavelength justification

added, flare science use case added Tritschler & Uitenbroek

02/14/2010

Major modifications to structure, update of

passband specifications, move of content to

DRD/PDD documents.

Wöger

02/24/2010 Minor changes, passband transmission requirements

follow from SNR requirements Wöger

03/25/2010

Updates to passband descriptions Ca II K, Hβ, Hα,

Ca II IR. Passband specification updated for Ca II

IR.

Wöger

05/07/2010 Proposed changes for an FPI Wöger

05/17/2010 Minor corrections, Appendix as suggested by

William McBride Wöger

08/12/2010 Clarifications & modifications. Release of Revision

G. Wöger

4/30/2012 Modifications in preparation for VBI red CDR Wöger

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SPEC-0054, Revision H Page iii

Table of Contents

1. INTRODUCTION ................................................................................................... 1

1.1 PURPOSE ............................................................................................................... 1 1.2 SCOPE ................................................................................................................... 1 1.3 DOCUMENT REVISIONS ............................................................................................ 1 1.4 APPLICABLE DOCUMENTS ........................................................................................ 1 2. MISSION ............................................................................................................... 2

3. EXAMPLE SCIENCE CASES ............................................................................... 3 3.1 MAGNETOCONVECTION ............................................................................................ 3 3.2 SUNSPOT STRUCTURE ............................................................................................ 3 3.3 MAGNETIC ELEMENT STRUCTURE AND DYNAMICS ..................................................... 4 3.4 CHROMOSPHERIC DYNAMICS ................................................................................... 4

3.5 FLARE DYNAMICS .................................................................................................... 4

4. INSTRUMENT PROPERTIES ............................................................................... 6 4.1 WAVELENGTHS AND PASSBAND WIDTHS ................................................................... 6 4.1.1 Ca II K line (393.4 nm) .................................................................................................... 8 4.1.2 G band (430.5 nm) ......................................................................................................... 9 4.1.3 Blue continuum (450.4 nm) ............................................................................................10 4.1.4 Hβ (486.1 nm) ...............................................................................................................11 4.1.5 Hα (656.3 nm) ...............................................................................................................12 4.1.6 TiO band (705.4 nm) .....................................................................................................13 4.1.7 Ca II IR triplet line (854.2 nm) ........................................................................................15 4.2 IMAGE QUALITY ..................................................................................................... 16 4.2.1 Spatial Sampling at the Diffraction Limit ........................................................................16 4.2.2 Adaptive Optics and Image Reconstruction ...................................................................16 5. REQUIREMENTS ............................................................................................... 17

5.1 SPECTRAL RANGE ................................................................................................. 17 5.2 FOV .................................................................................................................... 17 5.3 STATIC ABERRATIONS ........................................................................................... 17

5.4 SPATIAL SAMPLING ................................................................................................ 17 5.5 RELATIVE PHOTOMETRY ........................................................................................ 17

5.6 SYNCHRONIZATION BETWEEN CHANNELS ................................................................ 18 5.7 MULTI-WAVELENGTH CADENCE .............................................................................. 18 5.8 SIGNAL-TO-NOISE RATIO ....................................................................................... 18 5.9 LOCATION OF INSTRUMENT .................................................................................... 18

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SPEC-0054, Revision H Page 1 of 18

1. INTRODUCTION

1.1 PURPOSE

The ATST includes a suite of instruments designed in partnership with the ATST project in order

to realize the top level science requirements as described in the ATST SRD (SPEC-0001). As one

member of the first-light instrumentation suite, the Visible Broadband Imager (VBI) is designed

to achieve a subset of the top level science requirements, specifically to record images at the

highest possible spatial and temporal resolution of the ATST at a number of scientifically

important visible wavelengths. This document describes the science requirements that flow down

to the VBI design in order to achieve this subset of the over-all science goals of the ATST

project.

1.2 SCOPE

Section 2 of this document defines the specific mission of the VBI instrument as a member of the

ATST first-light instrumentation suite. In order to achieve this mission, the VBI must be capable

of providing the necessary data type, data quality, and data throughput to allow scientists to solve

crucial problems in solar research. These problems, and the roles that the VBI will play in

providing the data necessary to solve these problems, are explored in section 3 of this document.

Achieving the performance necessary to accomplish the science as described in section 3 places

severe requirements on the optical system of the VBI instrument which are analyzed and

described in Section 4. Section 5 is a summary of the requirements from section 4 that documents

the flow down of requirements from the ATST SRD.

1.3 DOCUMENT REVISIONS

Changes to this document have to be coordinated with the responsible author.

1.4 APPLICABLE DOCUMENTS

[1] ATST Science Requirements Document (SRD), SPEC-0001

[2] ATST Glossary and Acronym List, SPEC-0012

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2. MISSION

The mission of the VBI is to record images from the ATST telescope at the highest possible

spatial and temporal resolution at a number of specified wavelengths in the range from 390 nm to

860 nm. This will be accomplished with an optical design that preserves the Strehl-ratio of the

image provided by the telescope as well as possible under the constraint of scientific

requirements, and that has a high optical throughput at all considered wavelengths. In addition,

the VBI is required to have the ability to allow for image reconstruction for improving image

quality beyond what is provided by the telescope AO system. To maximize the FOV at the

required spatial sampling, the VBI must have large format arrays in its image plane.

The VBI will provide high-quality imaging through filters with relatively broad passbands to

optimize throughput. Its design will stress high cadence and short exposure times at the expense

of information in the spectral domain. With suitable detector QE, the wide bandpass and high

reflectance/transmission optics will allow exposure times short enough to effectively “freeze” the

atmospheric turbulence to apply speckle interferometric or deconvolution image reconstruction

techniques.

The wavelength regimes of the filters used in the VBI will be designed to span a range of

temperatures within the solar photosphere and chromosphere. It is required that the VBI be

capable of observing these two atmospheric regions simultaneously. In addition, the VBI will be

required to operate in a “multi-spectral” mode in which interleaved images in separate

wavelength regions are taken in rapid sequences to produce quasi-simultaneous “movies” of

different layers in the solar photosphere and chromosphere.

The VBI, or a subset of its components, is planned to be the “first light” instrument of the ATST.

It is therefore imperative that it must be completely tested and operational before the telescope

comes online. It will often be used as a “context” instrument, providing high quality imagery in

support of observations with other instruments, and should, therefore, have high availability and a

short preparation time.

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3. EXAMPLE SCIENCE CASES

In the following we concentrate on selected science cases.

3.1 MAGNETOCONVECTION

Studies of the solar convection zone and its radiative layer known as the photosphere are directly

relevant to understanding the structure and evolution of all Sun-like stars. In addition, the

mechanisms of magnetic field generation, evolution, and dispersion in the convection zone

remain unclear. The study of convective flows with magnetic fields, or “magnetoconvection”,

remains at the forefront of computational and observational astrophysics. The Sun is so far the

only star on which we can directly resolve the convection flows on the surface and thus observe

the interaction of the flows with magnetic field structures such as sunspots, pores, and the smaller

“magnetic elements” that make up active region and quiet-Sun network fields.

The primary objective of the VBI is to image the solar photospheric flow fields and magnetic

structures at the highest possible spatial and temporal resolution available from the ATST. By

compiling movies of the flow fields and magnetic structure interactions in several visible-light

wavelength bands, scientists can compare the observations to numerical simulations of

magnetoconvection in order to verify and validate the model results. Once validated by surface

observations, these magnetohydrodynamic (MHD) simulations can give scientists virtual

windows to the otherwise invisible physics of the solar interior convection zone. Although

developing techniques such as helioseismology can directly measure large-scale structures in the

interior, MHD simulations remain the only tool that can answer questions about the small-scale

turbulent structure of the convection, its relation to the solar rotational profile, and its interaction

with magnetic elements.

Although the VBI cannot measure magnetic fields directly (because it lacks both the spectral

resolution and polarization detection capabilities necessary), it will observe magnetic fields in the

photosphere through their radiative signatures in, for example, molecular band head regions of

the solar spectrum. These band heads are highly temperature sensitive and thus display magnetic

structures with high contrast relative to the non-magnetic surrounding photosphere.

3.2 SUNSPOT STRUCTURE

Sunspots are the largest and highest field strength magnetic structures visible in the photosphere.

Once thought to be fairly monolithic “pillars” of magnetic field surrounded by the filamentary

“penumbral” magnetic fibers, sunspots are now known to be highly structured formations with

convective elements inside of the umbra (“umbral dots”), varying magnetic fields across the

umbra, highly dynamic flows and variations of angles within the penumbra, and “light bridges”

which often demarcate fracture lines on which sunspots eventually break up. Particularly in the

case of the penumbra, scientists are still without a clear model of how sunspots form, evolve, and

disperse. The newly discovered “dark cores” of penumbral filaments also remain unexplained. As

with the magnetoconvective flows in the photosphere, some of the only clues about sunspot

structure have come from MHD numerical simulations. However these also require verification

and validation through detailed observations of sunspot structure and dynamics using the VBI on

the ATST.

The VBI will obtain high resolution, fast cadence, movies of penumbral fiber formation, umbral

flows, and interactions. It will use several wavelength bands to discriminate between various

layers in sunspot umbral and penumbral structure. As in the case of the magnetoconvective

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studies, direct measurement of the magnetic fields will not be possible, but highly accurate

structure and dynamics information will be obtained through non-polarimetric observations.

3.3 MAGNETIC ELEMENT STRUCTURE AND DYNAMICS

Magnetic elements are currently the smallest detectable magnetic structures in the solar

photosphere. They possess kilogauss field strengths but have observed diameters on the order of

only 100 km giving flux values on the order of 1016

Mx or less. Numerical MHD simulations

predict that these structures may have diameters as low as 10 km. MHD simulations also predict

that magnetic elements are concentrated in the down flow plumes of the intergranular lanes. The

dynamics of magnetic elements have been observed with 150 km resolution and found to be

dominated by fluid motions that result in constant merging and splitting of elements on time

scales of 100 seconds. The fluid motions are presumably those of the local intergranular lane

flows in which the magnetic elements are trapped. It has been hypothesized that magnetic

elements are channels for MHD and acoustic wave energy into the upper solar atmosphere,

although current observations cannot achieve high enough spatial or temporal resolution to verify

this. Whether magnetic elements are concentrations of “pre-existing” photospheric flux that was

originally generated in large-scale structures such as sunspots or whether they are generated

locally by fast dynamo action in the photosphere is also a key question in solar physics.

In addition, magnetic elements in active region plage are the source of “faculae”, bright structures

seen very near the solar limb in visible light wavelength ranges. Faculae are thought to be the

major source of the “excess” solar irradiance reaching Earth during periods of maximal sunspot

activity. However this aspect of facular brightness remains hypothetical and requires further

observations to verify. By observing faculae embedded in the granulation at the very high spatial

and temporal resolutions afforded by the ATST scientists will be able to investigate the detailed

contribution to the irradiation that faculae provide. Such observations require the highest possible

spatial resolution images at a minimum cadence of 10 seconds directly at different positions from

disk center towards the solar limb. Limb observations require very low scattered light from both

the ATST and the VBI.

3.4 CHROMOSPHERIC DYNAMICS

Convective processes in the solar atmosphere generate acoustic waves with different frequencies.

Those waves with frequencies below the acoustic cutoff frequency at about 5 mHz are reflected

in the gravitationally stratified solar atmosphere and are trapped below the photosphere. Higher

frequency waves are able to travel upwards into the chromosphere. According to numerical

simulations these waves form shocks in the low density chromosphere and contribute to its

heating. How much mechanical heating is contributed, where it is deposited and how much wave

energy is generated is still very uncertain. High spatial resolution observations at different

wavelengths that reflect atmospheric variations at different heights with a cadence of about 3

seconds or less are required to provide answers to these questions. VBI will be able to take

interleaved images in separate wavelength passbands in rapid sequences to produce quasi-

simultaneous “movies” of changes in the intensity from the solar photosphere and chromosphere.

3.5 FLARE DYNAMICS

It is commonly believed that solar flares represent a process of rapid transformation of the

magnetic energy of active regions into the kinetic energy of energetic particles and plasma flows,

and heat. An important goal for the ATST will be to study the small-scale processes in solar

flares. High-resolution hard X-ray and microwave observations have shown that flare bursts

occur in finely structured elementary bursts on time scales from tens of milliseconds to a few

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seconds. Potentially, these elementary bursts could be studied at higher spatial resolution in the

visible to obtain complimentary data on their physical properties. The VBI will be able to take

high cadence observations in chromospheric lines like Hα and Ca II 854.2 nm, and photospheric

continua simultaneously and at the highest resolution achievable with the ATST. This will allow

scientists to monitor the reaction of both the chromospheres and the photosphere to the fast

reorganization of the magnetic field and the corresponding heating phenomena at very small

spatial and short temporal scales. The VBI will also serve as the context instrument for spectro-

polarimetric measurements of flare regions, providing both higher cadence and a larger field of

view than the VISP, NIRSP and VTF.

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4. INSTRUMENT PROPERTIES

The main requirement of the VBI design is

1. to preserve the image quality (consistent with ATST’s diffraction limit error budget)

across the maximum FOV accessible under the constraint of scientific requirements, and

2. to deliver images at high temporal cadence.

The wavelengths for the VBI filters are specified with the purpose of providing sensitivity in the

wavelengths generated in altitudes from the deep photosphere into the chromosphere.

4.1 WAVELENGTHS AND PASSBAND WIDTHS

Observable wavelengths are divided into two channels, one in the red, and one in the blue. Each

of these two channels will have passbands targeted to chromospheric as well as photospheric

structure, so that the photosphere and chromosphere can always be monitored simultaneously by

using both channels in experiments investigating the connection between the photosphere and the

chromosphere. Given the fixed nature of the individual VBI passbands, the VBI targets the

observation of structure and dynamics from phenomenology at high temporal cadence and the

highest achievable spatial resolution, giving up spectral information in favor of high throughput

and simplicity of operation and data reduction. In chromospheric lines in particular, this

compromise implies that temperature and velocity information are inextricably mixed. The VBI

will therefore often be a survey instrument providing evidence of structure and dynamics at high

spatial scales and short temporal scales that can then be followed up with spectro(-polarimetric)

observations to provide better insight into the physics behind the fine structure.

The filter passbands of the VBI are chosen to be wide enough to provide excellent throughput, yet

narrow enough to isolate spectral features with a well-defined height resolution in the solar

atmosphere. The precise width and central wavelength of each passband are chosen, within the

above constraints, to optimize contrast as determined from numerical simulations or spectrally

resolved observations. A list of target spectral features and their approximate central wavelengths,

and formation heights are provided in Table 1 where the top four passbands are for the blue

channel and the bottom three for the red.

The filter passbands must be changeable within the constraints of the required temporal cadence

(< 3.2s per wavelength). Within each channel, filters must be changeable in any order, and at a

given time the filter selection mechanism should allow any combination of passbands between

different channels.

Due to the optical nature of passband filters, a shift of their CWL over the FOV is unavoidable

when the filters are used in a collimated beam. However, the drift in CWL and/or width over the

FOV cannot be so severe that the scientific goals addressed by the passband filter are

compromised. This would be the case if the drift exceeded the CWL tolerance given below or ½

the FWHM of the passband.

When passband filters are operated in a telecentric beam, both the FWHM and the Strehl-ratio are

impacted by the f-ratio (or f-number) of the beam incident on the filter; the f-number itself is

constrained by the availability of sufficiently large filters with adequate optical quality. This fact

constrains, in particular, the FOV.

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CWL [nm] Spect. Feature Height Purpose ATST SRD

393.4 Ca II K line Chromosphere Magnetic element morphology,

chromospheric dynamics, context

3.1,

3.2.1-3,

3.3.2

430.5 G band (CH) Photosphere Magnetic element morphology and

evolution, granular contrast, context

3.1.1-4,

3.1.8

450.4 blue continuum Photosphere Granular contrast, photospheric

dynamics

3.1.1-4,

3.1.8,

3.2.1-3

486.1 Hβ line Chromosphere Chromospheric morphology and

dynamics, context

3.1.1-4,

3.1.8,

3.2.1-3,

3.3.2

656.3 Hα line Chromosphere Chromospheric morphology and

dynamics, context

3.1.1-4,

3.1.8,

3.2.1-3,

3.3.2

668.4 red continuum Photosphere Granular contrast, photospheric

dynamics

3.1.1-4,

3.1.8,

3.2.1-3

705.4 TiO band Photosphere Sunspot umbral structure and

dynamics

3.1.1-4,

3.1.7-8

TBD

Table 1: VBI wavelength specifications. Top four wavelengths are for the blue channel, bottom four for the

red.

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4.1.1 Ca II K line (393.4 nm)

The K line filter should be narrow enough to ensure that chromospheric heights are probed, but

wide enough to allow for sufficient photon flux for exposures that still have sufficient signal-to-

noise ratio for speckle reconstruction, but are still short enough to freeze seeing. This is a

challenge at these blue wavelengths and may only be possible at sufficiently small airmass. The

passband should include the K2v emission feature to maximize chromospheric contribution and

sensitivity to both magnetic fields and acoustic shock signatures. The K line filter will also serve

to connect high resolution solar observations to stellar activity studies for which the K-line index

is commonly used magnetic activity proxy.

Property Value

CWL 393.36 ± 0.035 nm

FWHM ≤ 0.1 nm

Figure 1: Ca II K line core and filter curve (2 cavities). The diamonds demark the center of gravity position

of wavelength and intensity when the passband within its CWL tolerances is applied.

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4.1.2 G band (430.5 nm)

The G band filter serves to study the evolution of small scale magnetic features, which have

excellent contrast in this molecular band. These images will be imaged at the resolution limit of

0.02 arcseconds of the telescope owing to the relatively wide passband of 0.5 nm, which allows

very short exposure times, freezing seeing and optimally allowing for image reconstruction

techniques. The G band also serves as a context wavelength band as it is used at most solar

observatories and for comparison with other telescopes.

While the CWL of the G band filter needs to be well characterized, its tolerance may be relaxed

due to the fact that the CH molecular band is broad.

Property Value

CWL 430.5 ± 0.2 nm

FWHM 0.5 ± 0.15 nm

Figure 2: G-band spectrum and filter curve (2 cavities). The diamonds demark the center of gravity position

of wavelength and intensity when the passband within its CWL tolerances is applied.

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4.1.3 Blue continuum (450.4 nm)

The 450.4 nm wavelength band is located in a region with relatively few lines in the blue and

provides slightly higher granular contrast than the G band. As in the G band, the high throughput

from the relatively broad filter will allow very high resolution imagery of the convective

photosphere.

The tolerance for the blue continuum filter CWL can be relaxed as the inclusion of absorption

lines is unavoidable. However, the filter’s CWL must be very well characterized to be capable to

compare observations with MHD models.

Property Value

CWL 450.4 ± 0.2 nm

FWHM 0.4 ± 0.1 nm

Figure 3: Blue continuum spectrum and filter curve (2 cavities). The diamonds demark the center of gravity

position of wavelength and intensity when the passband within its CWL tolerances is applied.

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4.1.4 Hβ (486.1 nm)

The Hβ line shares its lower level with Hα line and has a similar chromospheric signature, but

with the benefit of slightly higher resolution at its shorter wavelength and higher spatial sampling

of the blue channel. It will provide a Hα like diagnostic in the blue channel that can be observed

with a photospheric passband in the red channel to sample photosphere and chromosphere strictly

simultaneously or together with Hα to provide chromospheric information at two different heights

to determine wave propagation properties. It is more suitable than the K line passband in this

respect because the Hβ line is more sensitive to velocities because of its steeper wings, like those

of Hα. The filter passband will be narrow and centered on the line core.

Property Value

CWL 486.136 ± 0.03 nm

FWHM ≤ 0.05 nm

Figure 4: Hβ line core and filter curve (1 cavity). The diamonds demark the center of gravity position of

wavelength and intensity when the passband within its CWL tolerances is applied.

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4.1.5 Hα (656.3 nm)

The hydrogen Balmer α line is the premier line to study chromospheric morphology and

dynamics. It has a 5 times higher opacity than Hβ and is therefore more sensitive to conditions in

the upper chromosphere. As for Hβ, the filter passband will be narrow and centered on the line

core. The Hα line shows remarkable variations in its shift, width and to a lesser degree in its

intensity levels, making the line exceptionally suitable for filtergraph observations showing large

contrast at a fixed wavelength, with the caveat that the signal in the passband is a mix of all three

effects which can therefore not be disentangled. In particular, the width of the Hα line is sensitive

to temperature because of the low atomic mass of hydrogen.

Property Value

CWL 656.282 ± 0.03 nm

FWHM ≤0.05 nm

Figure 5: Hα line core and filter curve (1 cavity). The diamonds demark the center of gravity position of


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4.1.6 Red Continuum (668.4 nm)

This spectral region of the red continuum passband is only very sparsely populated with spectral

lines, and the passband offers a good continuum image at lower spatial resolution and contrast

than the blue continuum passband because of its longer wavelength, but benefits from the larger

field of view of the red channel, and potentially more benign seeing conditions in the red.

Property Value

CWL 668.4 ± 0.2 nm

FWHM 0.4 ± 0.1 nm

Figure 6: Red continuum spectrum and filter curve (2 cavities). The diamonds demark the center of gravity

position of wavelength and intensity when the passband within its CWL tolerances is applied.

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4.1.7 TiO band (705.8 nm)

The opacity in the TiO band at 705.8 nm is highly temperature sensitive. In normal photospheric

conditions the band mostly vanishes, and the passband just provides a continuum image. In the

reduced temperatures of sunspot penumbrae, and umbrae in particular, the TiO line depths

increase significantly. This temperature sensitivity can therefore be used to good advantage to

bring out temperature fluctuations in the magneto-convective environment of sunspots which are

much reduced compared to those in normal granular convection.

The filter will have its peak transmission specified not exactly on the bandhead (located at

705.4 nm) but slightly towards the red part of the spectrum so that the measured intensity will

originate almost completely from within the band, and will only be minimally contaminated by

the nearby continuum.

Property Value

CWL 705.8 ± 0.2 nm

FWHM 0.6 ± 0.2 nm

Figure 7: TiO band and filter curve (2 cavities). The diamonds demark the center of gravity position of


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4.1.8 TBD

The specifics of this filter will be inserted, once a decision on the wavelength of this filter has

been made.

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4.2 IMAGE QUALITY

4.2.1 Spatial Sampling at the Diffraction Limit

In accordance with imaging and image reconstruction requirements, the VBI is required to have

at least two detector pixels spanning the diameter of a resolution element λ/D (in which a perfect

telescope concentrates approximately 50% of the energy of a point source) at each of its specified

wavelengths. Recording an image at this critical sampling in the Fraunhofer G-band (430.5 nm)

over up to 2′ × 2′ FOV requires 12k × 12k pixels, which is well beyond the dimensions of

currently available fast-readout CCD detectors. Nevertheless, the VBI will be required to have

optical fidelity over up to 2′ × 2′ FOV so that larger format detectors may be employed in the

future.

When critical sampling is implemented near the shortest VBI wavelength in one channel, the

image plane will be spatially oversampled at the longer wavelengths in the same channel if the

employed optical path is identical. This results in a different covered area on the Sun for each

wavelength in a channel, with common coverage determined by the smallest FOV.

4.2.2 Adaptive Optics and Image Reconstruction

The VBI will be capable of operating in conjunction with the adaptive optics system of the ATST

in order to achieve the highest possible spatial resolution of the telescope. This necessitates a

mounting in the rotating Coudé lab on one of the horizontal optical benches.

Post-facto image reconstruction techniques are needed to meet the requirement of diffraction

limited image quality over the entire FOV. Depending on the reconstruction technique employed,

the VBI will be required to receive WFCS data that are necessary to ensure a photometrically

correct reconstruction.

Modern and feasible reconstruction techniques include MFBD and speckle imaging. These

techniques require a sufficient signal-to-noise ratio to be capable of reconstructing signals up to

the diffraction limit, implying the need for high optical throughput.

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5. REQUIREMENTS

5.1 SPECTRAL RANGE

Requirement: 390 nm to 860 nm

Goal: 330 nm to 1.1 µm

Priority: 1

Source: SRD 2.8, coverage of photospheric and chromospheric features

5.2 FOV

Requirement: 2 arcminute square (Ca II K, G band, blue continuum),

2 arcminute circular (Hβ)

2 arcminute square (red continuum, TiO band, TBD),

2 arcminute circular (Hα)

Goal: 2 arcminute square (all wavelengths)

Priority: 1

Source: SRD: 3.2.4, 3.3.1, 3.3.2

5.3 STATIC ABERRATIONS

Requirement: known static aberrations with p-v amplitude < λ/2 for aberrations of higher order

than focus, at all observed wavelengths

Goal: known static aberrations with p-v amplitude < λ/4 for aberrations of higher order

than focus, at all observed wavelengths

Priority: 1

Source: SRD: 3.1.1, 3.1.2, 3.1.3, 3.1.4, 3.1.5, 3.1.7, 3.2.4, 3.3.2

Note: Fluctuating, non-common path aberrations are removed by post-facto image

reconstruction algorithms. Some reconstruction algorithms, such as speckle

reconstruction algorithms, are not capable of removing the effect of static

aberrations. However, if the static aberrations are well characterized, they can be

removed within the above specified limits by an extra processing step.

5.4 SPATIAL SAMPLING

Requirement: Nyquist-Shannon sampling at 430.5 nm (VBI blue channel)

Nyquist-Shannon sampling at 656.3 nm (VBI red channel)

Goal:

Priority: 1

Source: SRD: 3.1.1, 3.1.2, 3.1.3, 3.1.4, 3.1.5, 3.1.7, 3.2.4, 3.3.2

5.5 RELATIVE PHOTOMETRY

Requirement: 2 × 10-2

I0 in saved image

Goal: 1 × 10-2

I0 in saved image

Priority: 1

Source: SRD: 3.2.2, 3.2.3, Sunspot umbra observations, granular contrast.

Remarks: Impacted by scattered & stray light, passband out-of-band rejection, camera

synchronization with modulation, image reconstruction algorithms

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5.6 SYNCHRONIZATION BETWEEN CHANNELS

Requirement: 10 ms

Goal: < 10 ms

Priority: 2

Source: SRD: 3.1.7, 3.1.8 (Alfvén wave propagation)

5.7 MULTI-WAVELENGTH CADENCE

Requirement: 3.2 seconds for one wavelength, including filter change, vibration settle and data

acquisition time for 80 frames.

Goal: 3 seconds

Priority: 1

Source: SRD: 3.1.1, 3.1.2, 3.1.4

5.8 SIGNAL-TO-NOISE RATIO

Requirement: ≈ 100 per delivered image

Goal: > 100

Priority: 1

Source: SRD 2.1: image reconstruction (achieve full correction over the entire FOV)

5.9 LOCATION OF INSTRUMENT

Requirement: Post-adaptive optics

Goal: NA

Priority: 1

Source: SRD 2.1: image reconstruction (provide maximum possible spatial resolution

with AO correction for post-facto image reconstruction algorithms)

Visible Broadband Imager (VBI) Instrument Science Requirements

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Transcript of Visible Broadband Imager (VBI) Instrument Science Requirements