Observational Cosmology Stephen Serjeant Pdf Converter

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. 623 Downloads. Abstract Half of the energy ever emitted by stars and accreting objects comes to us in the far-infrared (FIR) waveband and has yet to be properly explored. We propose a powerful Far-InfraRed Interferometer mission, FIRI, to carry out high-resolution imaging spectroscopy in the FIR. This key observational capability is essential to reveal how gas and dust evolve into stars and planets, how the first luminous objects in the Universe ignited, how galaxies formed, and when super-massive black holes grew. FIRI will disentangle the cosmic histories of star formation and accretion onto black holes and will trace the assembly and evolution of quiescent galaxies like our Milky Way.

Perhaps most importantly, FIRI will observe all stages of planetary system formation and recognise the birth of planets via its ability to image the dust structures in planetary systems. FIRI is an observatory-class mission concept: three cold, 3.5-m apertures, orbiting a beam-combining module, with separations of up to 1 km, free-flying or tethered, operating between 25 and 385 μm, using the interferometric direct-detection technique to ensure μJy sensitivity and 0.02” resolution at 100 μm, across an arcmin 2 instantaneous field of view, with a spectral resolution, R 5,000 and a heterodyne system with R 1 million.

Although FIRI is an ambitious mission, we note that FIR interferometry is appreciably less demanding than at shorter wavelengths. Half of the energy ever emitted by stars and accreting objects comes to us in the far-infrared (FIR) waveband and has yet to be properly explored.

We propose a powerful Far-InfraRed Interferometer mission, FIRI, to carry out high-resolution imaging spectroscopy in the FIR. This key observational capability is essential to reveal how gas and dust evolve into stars and planets, how the first luminous objects in the Universe ignited, how galaxies formed, and when super-massive black holes grew.

FIRI will disentangle the cosmic histories of star formation and accretion onto black holes and will trace the assembly and evolution of quiescent galaxies like our Milky Way. Perhaps most importantly, FIRI will observe all stages of planetary system formation and recognise the birth of planets via its ability to image the dust structures in planetary systems. It will thus address directly questions fundamental to our understanding of how the Universe has developed and evolved—the very questions posed by ESA’s Cosmic Vision. What are the conditions for stars to form, and where do they form?.

How do stars evolve as a function of their interstellar environment?. In which conditions do planets form around stars?.

Stephen Serjeant

How were the first luminous objects in the Universe ignited? How did the first stars form and evolve?. How did the history of stars and supernovae give rise to current chemical element abundances?.

What is the history of super-massive black holes and how do they interact with their host galaxy?. What is the nature of the FIR background, and of early, deeply embedded star formation? The FIR region of the electro-magnetic spectrum is the last major band where poor angular resolution and lack of sensitivity hinders progress.

Pathfinders Herschel and SPICA (see ) will provide major advances in sensitivity, but will lack the angular resolution necessary to resolve the cosmic FIR background radiation or to undertake detailed studies of individual objects. ALMA and the James Webb Space Telescope ( JWST) will provide high angular resolution and sensitivity at shorter and longer wavelengths, but the crucial band between 25 and 300 μm is not covered by any comparable instrument: there exists a crippling lack of observational capability in the FIR, despite the vital role this band plays in exploring the formation and evolution of Active Galactic Nuclei (AGN), galaxies, stars and planetary systems, and the development of life-sustaining environments. Fig. 1 With the advent of ALMA and VLTI, the FIR gap is deepening (picture courtesy of T. Wilson) FIRI will make unique and key contributions to our understanding of the Universe, near and far. It will peer through the dust that shrouds stellar nurseries in our Galaxy, de-mystifying the process by which stars and planets are born. It will image proto-stellar and debris disks at the peak of their spectral energy distributions (SEDs), where the brightness is 1,000× that at a wavelength of 1 mm, with exquisite spectral resolution and sensitivity, revealing how planetary systems form out of gas, dust and ice.

While the Hubble Space Telescope ( HST) has produced beautiful pictures of merging and star-forming galaxies, crude submm observations have shown that the real action is in the FIR. FIRI’s angular resolution will break through the confusion limit and allow us to determine the properties and internal structure of distant star-forming galaxies, and to examine the enigmatic symbiosis between host galaxies and their AGN. The earliest metal-poor galaxies will be very highly redshifted (i.e. At large look-back times), and we may even be able to detect their formation via molecular hydrogen emission red-shifted into the FIR. This would provide a unique probe of first light—the formation of the earliest stars in the Universe and the ensuing re-ionisation of the Universe. Here, we outline the FIRI observatory-class mission concept: three cold, 3.5-m apertures, orbiting a beam-combining module, with separation of up to 1 km, free-flying or tethered, operating between 25 and 385 μm, using the interferometric direct-detection technique to ensure μJy sensitivity and 0.02” resolution at 100 μm, across an arcmin 2 instantaneous field of view, with a spectral resolution, R 5,000 and a heterodyne system with R 10 6.

Although FIRI is an ambitious mission, we note that FIR interferometry is appreciably less demanding than at shorter wavelengths (Fig. The untapped potential of FIR astronomy is most clearly illustrated by considering the three main components that dominate the electromagnetic energy content of the Universe (Fig. The dominant component is the microwave background produced by the primordial Universe at recombination ( z 1,089). The second most important is the FIR background, produced by galaxies in the young Universe.

The third is the optical background dominated by evolved stars/galaxies and AGN. The first and third of these components have now been mapped in detail over the entire sky, while virtually no sky has been imaged in the FIR to any reasonable depth. Fig. 4 SMA image of the dust continuum emission from the transitional disk around LkH α 330, a G3 pre-main sequence star in Perseus.

The putative outer gap radius of 50 AU, derived from an SED analysis, is confirmed. Gas is detected well inside the transitional radius, raising the possibility that the gap is caused by a planetary companion The FIR wavelength region is thus the least explored part of the electro-magnetic spectrum yet it provides uniquely powerful tools to study material associated with the earliest evolutionary stages of galaxies, stars and planets. This waveband allows us to directly probe objects during their formation. We gain access to unique science—the peak of the FIR background radiation and emission from cold, proto-stellar cores.

A large, single, actively cooled aperture in space would provide unprecedented increases in sensitivity and mapping speed over any existing or planned facility. However, it is clear that to make significant advances on Herschel and SPICA in all areas of astronomy, from nearby planetary systems to the highest redshift galaxies, even the largest conceivable dish would be inadequate. These areas call for exquisite angular resolution, around 0.02 arcsec at 100 μm, as well as sufficient sensitivity to allow photon-starved spectroscopy across an arcmin 2 field of view in the 25- to 385-μm band. These science requirements lead us inexorably to an interferometer.

3 The FIRI science programme. A space-based FIR interferometer will enable the following unique science: 3.1 Planets and life The FIR is a critically important wavelength range for studying the origin and evolution of planetary systems. Most solar systems, including our own, are pervaded by dust, which is very bright at FIR wavelengths.

By studying the structure and dynamics of this dust, with many times the effective sensitivity of ALMA since we are operating at the SED peak (for 10–120-K dust), we gain information on how such systems were formed. More importantly, we can infer the presence and orbital motions of planets, which influence the distribution of the dust. Complementary to this, the FIR is the natural place to observe the emission of organic molecules through rotational lines or low-lying vibrational bands that could literally be the building blocks of life in the Universe, including water which has a profound influence on both the formation of the planets as well as for life (e.g.

By better understanding the contents and chemistry of the interstellar medium (ISM) throughout the wide range of different environments found in the Milky Way and nearby galaxies, it will be possible to get a better idea about the potential for life in these planetary systems. 3.2 The interstellar medium (ISM) in quiescent galaxies In relatively quiescent local galaxies, including the Milky Way, the bulk of the ISM is relatively cool, with characteristic temperatures of order a few tens to a few 100 K. Therefore, the chemistry and dynamics of the ISM in these galaxies is often accessible only through FIR observations, as “cool” material radiates most strongly in the FIR. In particular, emission lines of CO offer a powerful tool for investigating dynamic processes in local (and distant) galaxies, and when combined with emission lines from atoms and ions of O, N and C, can be used to build up a picture of the chemical composition of the ISM. Other useful probes of the “cool” ISM include emission and absorption features from solid grains, most notably polycyclic aromatic hydrocarbons (PAHs), and large amorphous silicates.

3.3 The ISM in active galaxies Conversely, the most luminous sources of radiation in galaxies are either very hot, very young stars, or the accretion disk around the central super-massive black hole (SMBH). Both hot young stars and accretion disks emit nearly all their energy at UV, X-ray and even gamma-ray wavelengths. However, these sources are in most cases embedded in large amounts of gas and dust; for example, the earliest stages of star formation invariably occur deep inside clouds of interstellar gas and dust—their raw material, and AGN require the accretion of large amounts of gaseous fuel onto the central SMBH. This surrounding gas and dust absorbs most or all of the UV and soft X-ray radiation which is directly emitted by the sources, and re-radiates the bulk of it in the FIR; the surrounding clouds are transparent at these wavelengths and the grains can achieve an energy balance between radiation absorbed at short wavelengths and emitted at long wavelengths. As a result, most of the energy originally emitted by the high-temperature primary sources is converted into FIR and submm radiation. Only by observing at these wavelengths can we measure fundamentally important parameters such as total energy budgets ( L bol), accretion disk geometries and star-formation rates. Indeed, there is already strong evidence that the growth of the SMBHs found in the centres of nearby galaxies probably takes place at the same time as the formation of the bulk of stars in the central bulge (e.g.

As a result, this growth likely adds to the energy output of a galaxy in the FIR, but not at any other waveband. FIR observations thus provide the only way to detect this SMBH growth; even hard X-ray observations cannot penetrate the Compton-thick shrouds of gas in a forming galaxy (e.g. 3.4 Cosmology. Fig. 5 SCUBA imaging of the 1-arcmin 2 field around an absorbed z = 1.8 QSO revealing a remarkable 400-kpc-long chain of galaxies, each with an obscured star-formation rate sufficiently high to build a massive spheroid in.

In summary: FIRI will thus revolutionise our knowledge of the formation of galaxies, stars and planetary systems and the development of life-sustaining environments. We will be able to probe the universality of the IMF across a range of galaxy environments and map out star- and planet-forming disks in stellar nurseries through resolved spectral lines. FIRI will break the cosmic background radiation into its constituent parts—many thousands of faint, dusty high-redshift galaxies, resolving them individually to yield otherwise hopelessly obscured information about their formation and evolution. It will root out Compton-thick AGN and differentiate between gas heated by active nuclei and stars, thus disentangling the formation histories of SMBHs and stars.

FIRI’s “discovery potential” is also extremely high. The biggest surprise would be if there were no surprises. FIRI is therefore capable of answering many of the most important questions posed by Cosmic Vision (Fig.

Fig. 6 Spitzer-IRS observations of NGC 6240. The signatures of AGN activity are apparent in this Compton-thick source from the identification of Ne V at 14.3 μm. These lines are redshifted to 40 μm for z  2, i.e. Beyond the coverage of JWST, but far short of ALMA’s wavebands. The luminosity distance of NGC 6240 is 100 Mpc, two orders of magnitude closer than z 2 AGNs, indicating the need for a sensitive FIR imaging spectrometer to identify Compton-thick AGN in the distant Universe 4 Summary of science requirements. FIR community workshops were held in Madrid (2004), Leiden (2005) and Obergurgl (2006). The participants have quantified the requirements of 26 use cases, resulting in the science requirements summarised in Table.

The most consistent requirement—found in nearly every Galactic use case, in many extragalactic cases and in virtually all of the most important use cases—is the need for high angular resolution, 0.02 to l arcsec. Sensitivity is an equally demanding driver, especially for extragalactic science. Requirements on spectral resolution span the whole range, up to several 10 6; extragalactic science and around half of the Galactic use cases are satisfied with R 5,000. We list a short title; the wavelength to be observed; the wavelength range; frequency or velocity range; the spectral resolution required; the area to be covered; the angular resolution required, the number of fields to be measured in order to fulfil the use case; the noise level to be reached; particular constraints on the interferometer and the dynamic range required. This table also shows the versatility of FIRI, although we don’t expect all the use cases to be completed in a 5-year period 5 Mission profile. 5.1 From science to mission requirements From the science requirements the following conclusions can be drawn: the mission should be able to resolve proto-planetary disks and distant galaxies down to a few AU at about 100 pc distance, or 100 pc at z 10, i.e.

Angular resolution of 0.02”. In terms of wavelength, the mission should provide access to the 28-μm rotational transition of H 2 and overlap with ALMA, so 25 to 385 μm. Most of the science requires extremely high sensitivity, so adequately large and cold telescopes are a necessity, as well as next-generation detectors. A field of view rather larger than the primary beam is also desirable for virtually all the high-profile science.

To achieve the required angular resolution, long baselines are necessary, 1 km, longer than can be achieved with rigid structures. A free-flying or tethered interferometer capable of delivering excellent image fidelity is thus the instrument of choice. Unlike the optical and radio regimes, where point sources are common, the FIR wavelength region is full of structure on every spatial scale, at the angular resolution of FIRI.

This has profound implications for the way data must be obtained. An interferometer works well only when good ( u, v) coverage is achieved, with attention to the short or zero spacings to allow full image reconstruction. In recent years, four studies of FIR/submm interferometers have been performed. These were the American SPECS and SPIRIT studies and the European ESA-FIRI and ESPRIT studies. Studies for the direct-detection interferometer SPECS (two cold 4-m telescopes plus central beam combiner) and the heterodyne interferometer ESPRIT (six ambient antennas) show that the incoherent and coherent variants of FIRI can take very different shapes. Neither concept on its own is capable of fulfilling all science requirements nor is any ready for immediate implementation. The science community’s demands on spectral resolution range from 5 to 10 6.

Observational Cosmology Stephen Serjeant Pdf Converter Full

The latter is only achievable with a heterodyne system, while greater sensitivity is available at lower spectral resolution (. 6.1 Direct detection interferometer A direct detection “optical” interferometer can be used to obtain high-angular-resolution images and spectra simultaneously over a wide field of view with a single instrument. Such an interferometer is analogous to the VLTI. For FIRI we envisage an extension of the optical interferometry techniques commonly used along the direction suggested by Mariotti and Ridgway , who described the possibility of combining an imaging interferometer with a Fourier transform spectrometer. A further natural extension of this “double Fourier” technique involves the substitution of a detector array for the single detector used in a traditional Michelson beam combiner, a spatial multiplexing method which enables the simultaneous observation of many contiguous primary beams to cover a wide FoV, and which is currently under development in the laboratory ,.

At FIR wavelengths, space-based optical interferometry is not as difficult as it may seem. JWST, currently in development for operation at 10× shorter wavelengths, may be the largest practical single-aperture space telescope.

Its 6.6-m primary mirror is comprised of 18 segments whose light must reach a common focus. That will be accomplished with a wavefront sensing and control system , a system now considered sufficiently mature for flight application. A direct detection interferometer is the optical analogue of a pair, or multiple pairs, of JWST mirror segments. Because wavefront sensing and control becomes easier with increasing wavelength, a FIR interferometer will be able to adopt proven JWST technology.

6.1.1 Direct detection interferometer principles. Fourier transform spectroscopy is analogous to the method employed famously by. In this case light from a source is split into two beams and recombined after a moving mirror inserts a variable time delay between one beam and the other.

The combined light can be recorded on a single detector or camera pixel. The output signal intensity, when plotted against the optical delay, makes an “interferogram.” The spectrum of the source can be constructed from the Fourier transform of the interferogram. If the optical delay range Rλ is sampled, the resulting spectral resolution will be R.

In imaging interferometry, spatial coherence (interference “fringe visibility”) measurements at many interferometric baselines (telescope spacings and orientations) are Fourier transformed to produce an image. Each baseline samples a single spatial Fourier component of the brightness distribution of the target scene, a component commonly identified by its coordinates in the spatial frequency ( u, v) plane. The Van Cittert–Zernike theorem explains that an image can be constructed without information loss if enough ( u, v) plane positions, or baselines, are sampled.

By combining these techniques, as suggested above, FIRI will do integral field spectroscopy. When equipped with a detector array of modest pixel count the FIRI field of view can readily be expanded to 1 arcmin 2, as illustrated in (Fig. Fig. 7 Interference fringes from field angles outside the primary beam ( red ellipse) can be recorded simultaneously in the separate pixels of a detector array. If light from a source located on the optical axis of the interferometer ( solid lines) is focused onto pixel ( x o, y o), then light from an off-axis source ( dashed lines) might reach pixel ( x o + δx, y o + δy) after traversing opposite arms of the interferometer.

The white light fringe packet in the interferogram from the latter pixel ( top) is displaced relative to the interferogram from the on-axis source ( bottom). The FoV is dictated by the number of pixels in the array 6.1.2 Description and key characteristics Michelson stellar interferometers can be constructed with any number of light collectors N, and the light may be non-redundantly combined pairwise in up to N( N – 1) / 2 baselines.

For FIRI, we restrict ourselves to one of the simplest configurations, three light collectors (three baselines) and a beam combiner, all of which are on different spacecraft. The entire array rotates in a plane perpendicular to the line of sight, keeping the optical pathlength differences external to the instrument to a minimum (a few cm). The residual pathlength difference will be well within the range (1 m) of the scanning optical delay line required both for spectroscopy and to equalise path lengths for off-axis field angles.

Guide stars are used to orient the spacecraft in absolute coordinates, and distinct sources in the science field of view—possibly NIR point sources—serve as phase references for the optical path external to the instrument. To synthesize as complete an aperture as possible, the baseline length B between the two telescopes can be varied continuously between 10 m and 1 km. The minimum baseline length is determined by how close the light collectors can approach each other, and depends primarily on the sizes of the sun shields, and the architecture of the spacecraft formation. The maximum baseline is dictated by the science requirements on angular resolution (0.02 arcsec). A direct detection interferometer has three distinct positive attributes. First, with sufficiently cold (4 K) optical elements and low-noise detectors, its sensitivity will be limited only by statistical noise in the sky background radiation. Second, a FoV wider than the primary beam can be observed without repointing the interferometer and resampling the ( u, v) plane.

Third, the interferometer can provide uninterrupted access to a wide range of wavelengths. Because the detectors typically operate well over a single octave, FIRI will use four different detector arrays to cover the range from 25 to 385 μm. Potentially, filters could be used to narrow the bandwidth and limit background photon noise when the source of interest is seen against a bright background.

A single mechanism can be used to provide the optical delay scan for all four spectral channels. 6.1.3 Optical delay lines Cryogenic optical delay lines been used in space before, notably in the Cosmic Background Explorer FIRAS instrument and the Cassini CIRS instrument. The FIRAS mechanism is comparable in many respects (operating temperature, lifetime requirement, physical size, and rate of motion) to the delay line mechanism needed for FIRI.

Darwin has similar requirements. 6.1.4 Beam combiners The beam-combiner is the heart of the interferometer system. FIR beams are more easily combined than mid-IR or shorter wavelength beams because of the relaxed tolerances on wavefront flatness. Metal mesh filters with customized spectral response functions are well suited to serve as beam combiners and dichroics to divide the broad FIRI wavelength range into octaves specified to permit an optimal match to the detectors. The beam combining instrument grows geometrically in complexity when additional light collectors are added to a direct detection interferometer, as an additional delay line and two more detector arrays, one for each Michelson output port, are needed for each interferometric baseline.

A three-telescope system should be feasible if the telescopes and beam combiner are launched together. To minimize complexity, a two-telescope system was considered in the SPECS study. We recommend further studies of two and three-telescope architectures. Considerable experience in the design and operation of imaging FTS systems in the FIR is now being established. Herschel/SPIRE and SCUBA-2 will both have an imaging FTS and will provide experience in spectrometer operation and data reduction. This will be extended by SPICA/ESI, for which an imaging FTS with transition edge sensor (TES) arrays is proposed. 6.1.5 Detectors In order to meet the performance requirements quoted in Section, multi-pixel detector arrays with NEPs of 10 − 20 W Hz − 1/2 will be needed.

The FIRI detectors will be based on superconducting sensors and will require an operating temperature of 50 mK. Superconducting TES arrays operating at 0.1 K are being implemented for SCUBA-2 at the James Clerk Maxwell Telescope and are planned for other ground-based submm telescopes. While these instruments are developing and demonstrating much of the fabrication and multiplexing technology appropriate for FIRI, they are optimised for higher backgrounds, requiring NEPs of 10 − 17 W Hz − 1/2. However, low-background TES arrays are also being developed for the proposed ESI instrument on the Japanese SPICA satellite (the subject of a complementary Cosmic Vision proposal). Cardiff University, SRON, and other SPICA/SAFARI collaborators are developing TES arrays with frequency domain multiplexing, with a target NEP of 10 − 19 W Hz − 1/2, deemed appropriate for SPICA, requiring a detector technology selection in 2009.

The implementation of SPICA/SAFARI will constitute a thorough development programme for TES array and systems technology, and sub-100-mK cooler running as the last element of a cryogen-free cooling system. A natural extension of the SPICA detector programme would be to develop detectors with NEPs of 10 − 20 W Hz − 1/2. Similar FIR detector programmes are underway in North America. There is also a great deal of overlap between the fabrication techniques, cooling needs, and basic sensitivity requirements of detectors suitable for the next generation X-ray mission ( XEUS or its equivalent), and a post- Planck CMB anisotropy mission (such as B-Pol). The largest FIRI detector arrays, for the shortest wavelength channel, will have 32 × 32 pixels if the primary beam is Nyquist sampled and the FoV is 1 arcmin 2, or 16 × 16 pixels if Nyquist sampling is determined unnecessary. Smaller array dimensions cover the FoV in longer wavelength channels because the primary beam is larger. 6.1.6 Cryo-coolers CEA-SBT is currently carrying out a strategic development programme for sub-100-mK cooler systems for future space science experiments, including XEUS and SPICA, and we expect these will be considered a mature technology in the next decade.

A hybrid 3He sorption cooler/Adiabatic Demagnetisation Refrigerator is under development for SPICA/SAFARI, which can provide continuous cooling at or below 75 mK, and NASA Goddard has already achieved cooling to 30 mK with a Continuous ADR. Telescope cooling to 4 K will demand a somewhat more powerful cryo-cooler than the cooler developed for JWST/MIRI.

The JWST cryo-cooler was recently declared to have matured to TRL 6, and it will be adopted for flight. Experienced cryogenic engineers say that the development path from the JWST cooler to a cooler powerful enough for FIRI is straightforward. 6.1.7 Operation of the interferometer A FIRI observation sequence comprises a slew to the target field, target acquisition (including lock on angle and zero path difference tracking), and science data acquisition.

The telescopes can move during the delay line scan as long as their positions are known accurately. Data analysis will be simpler if the delay line is scanned faster than the time it takes for a telescope to move a distance equal to its diameter. The detectors will be calibrated at regular intervals.

The measurements continue until all of the desired baselines are sampled. A direct detection interferometer has a great deal of flexibility in its operation, enabling the FIRI to satisfy optimally a variety of measurement requirements. For example, if a particular observation calls for angular resolution coarser than 0.02 arcsec, spectral resolution. Detector array development to achieve very low NEPs and high read-out rates;. Thermal model validation and verification. What is the best optical design and how do we do stray-light suppression?. Beam combiner engineering design.

High specific impulse propulsion system. Power system to drive thrusters, operate mechanisms, detectors, cryo-coolers, communication system, etc.

How best to actuate mechanisms intended to work at 4 K with minimal parasitic heat load?. Sky density and quality of phase reference sources in the field of view?. What is the best AIV approach?

6.2 Heterodyne interferometer The heterodyne interferometer will be the instrument of choice when high spectral resolution is required. It provides this resolution naturally in the mixing and down-converting process; at the same time, unlimited amplification and splitting of the down-converted (Intermediate Frequency, IF) signal is possible.

A description of the principle is given below, followed by a more detailed description of key elements: mixers and local oscillators; correlators; and observing modes, metrology and ( u, v) plane filling. 6.2.1 Heterodyne interferometer principles In a heterodyne system the sky signal is combined with an internally generated Local Oscillator (LO) signal. When this combination is done in a non-linear mixing element, the result is a copy of the sky signal but amplified and down-converted from the THz range into the GHz range. In this regime, low-noise amplifiers are available to further amplify the signal. At this stage the signal is electrical rather than optical. The amplified signal can be split when needed and correlated with other “sky” (or better, amplified IF) signals in digital correlators.

The only actively cooled parts in the spacecraft at this time are the detectors (mixers). The high angular resolution guarantees that the telescope background is not an issue. An ambient telescope, like Herschel, suffices at these early stages. There are several ways of correlating the IF signals.

The ALMA case is based on a XF correlation scheme, while the advantage of novel FFT Spectrometers is that they can operate in a FX correlation scheme. For FIRI, a hybrid correlator (FX then XF) may well be the best option. A big advantage of a heterodyne interferometer is that the very high spectral resolution guarantees a very high correlation length; delays can thus be applied electronically or in software based on a geometrical model. This feature makes it possible to keep the metrology of a multi-telescope system rather relaxed, i.e.

Observing is possible while the telescopes are moving. In fact, correlation is possible with a prediction of the position of the spacecraft within tens of cm while the metrology of spacecraft position with respect to each other should be a few μm to guarantee good phase calibration. The heterodyne concept could make use of a central correlator, like the direct-detection beam combiner, but a distributed correlator (in all available spacecraft) is also possible and probably preferred. Since moving the telescopes while observing is not a problem, ( u, v) plane sampling can be achieved by letting the dishes move outwards (after acceleration), sampling during this floating phase. Deceleration to a stop and acceleration inwards is necessary to bring the array back to its initial configuration, after which the array can move to another source.

Since delays can be done electronically or in software, the array could even fly in a 3-D configuration, thereby mitigating collision risks. The heterodyne interferometer thus consists of several telescope spacecraft with, in their focal planes, several heterodyne mixers tunable at spot frequencies in the FIR. Mixers receive their LO signal from dedicated Local Oscillators, with phases tuned to a master LO phase distributor which distributes phases using a detailed geometrical model.

These signals are down-converted in the mixers and amplified several times. The IF signals (which are led to the service module of each spacecraft) can be split and distributed to the other spacecraft for correlation, where delays will be applied in electronics and software, without moving optical parts.

Observational Cosmology Stephen Serjeant Pdf Converter Youtube

An alternative would be to send all IF signals to a central correlator where visibilities are calculated and sent to Earth at a low data rate. Mosaicing will be possible during each observation, increasing the field of view of the heterodyne FIRI complement. 6.2.2 Mixers and local oscillators Based on knowledge of ground-based telescopes and the developments for Herschel/HIFI, we have identified two possible types of mixer: Hot Electron Bolometer (HEB; 1.5 THz) and Superconductor–Insulator–Superconductor (SIS.