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To properly operate a temporal hyvpertelescope. the optical path mocdulators are driven by an 8 channel function e&enerator and related. high. voltage electronics (not drawn in the picture). | To properly operate a temporal hypertelescope, the optical path modulators are driven by an 8 channel function generator and related high voltage electronics (not drawn in the picture). |
The output voltage drives the optical path mocdulators with a full span in the range of tens yam and a typical nanometric sensitivity. | The output voltage drives the optical path modulators with a full span in the range of tens $\mu m$ and a typical nanometric sensitivity. |
Phe electronic gain and the voltage generator slopes allow to set the v7; frequencies at the proper values. | The electronic gain and the voltage generator slopes allow to set the $\nu _i$ frequencies at the proper values. |
In such a configuration. we can theoretically eet the same imagine properties as for the first. classical design using spatial pupil densification. | In such a configuration, we can theoretically get the same imaging properties as for the first classical design using spatial pupil densification. |
Lhe breadboard described in this paper and the related experimental results reported in the next paragraphs aim to demonstrate the validity of this new concept. | The breadboard described in this paper and the related experimental results reported in the next paragraphs aim to demonstrate the validity of this new concept. |
Our experimental set-up. (see figure 5)) has been designed and implemented thanks to the different skills developed in our team for two decades (Allemanetal.1995: Simohamect&Revnaucd199 llussctal 2001: Perrinctal 2006:: Olivieretal 2007)). | Our experimental set-up (see figure \ref{schema_complet}) ) has been designed and implemented thanks to the different skills developed in our team for two decades \citealt{All}; \citealt{S}; \citealt{H}; \citealt{Perrin}; \citealt{O}) ). |
Consequently. our ΕΙ) experimental test bench uses optical fibres ancl couplers for the cilferent optical functions to be implemented. | Consequently, our THT experimental test bench uses optical fibres and couplers for the different optical functions to be implemented. |
Llowever. we would like to stress that the use of guided. optics components is not mandatorv for the implementation of a THT. | However, we would like to stress that the use of guided optics components is not mandatory for the implementation of a THT. |
A ‘lassical design with classical components could be chosen if preferred. | A classical design with classical components could be chosen if preferred. |
“Phis point will remain a minor one as long as we will focus more on the demonstration of TIET. principle mn on the technological aspects. | This point will remain a minor one as long as we will focus more on the demonstration of THT principle than on the technological aspects. |
The following items give 1e &eneral framework of our experimental study. | The following items give the general framework of our experimental study. |
The following sections summarize the “VIP bench structure. | The following sections summarize the THT bench structure. |
It consists of three main parts (cf Fig.5)): a star simulator. a telescope array and a combining interferometer. | It consists of three main parts (cf \ref{schema_complet}) ): a star simulator, a telescope array and a combining interferometer. |
The calibrated: object is the first. subsystem: required. for testing the imaging capability of a ΕΕ. | The calibrated object is the first subsystem required for testing the imaging capability of a THT. |
For this. first experimental demonstration. the selected astronomical target is a binary star with a convenient angular separation ancl adjustable For this purpose. the object consists of two tips of monomode Panda fibres glued on a V-groove. | For this first experimental demonstration, the selected astronomical target is a binary star with a convenient angular separation and adjustable For this purpose, the object consists of two tips of monomode Panda fibres glued on a V-groove. |
These monomode waveguides are fed by two independent Distributed FeedBack lasers (DEB) with the same emitting wavelength. and. act as two incoherent point like sources. | These monomode waveguides are fed by two independent Distributed FeedBack lasers (DFB) with the same emitting wavelength and act as two incoherent point like sources. |
This way the object is spatially incoherent and the dyvnanics is controlled. by adjusting the laser driving currents. | This way the object is spatially incoherent and the dynamics is controlled by adjusting the laser driving currents. |
A set of doublets ancl collimating lenses allows to. provide an angular intensity distribution compatible with the spatial frequencies ov sampled by our telescope array. | A set of doublets and collimating lenses allows to provide an angular intensity distribution compatible with the spatial frequencies $u$ sampled by our telescope array. |
In our experiment the angular separation (5. as seen by the telescope array. is 23.75pad. | In our experiment the angular separation $\theta_{0} $, as seen by the telescope array, is $23.75 \mu rad$. |
As our instrument ts designed for a linear input polarization. a polarizing cube is inserted in the doublet spacing in order to select. and fixes a linear vertical input polarization (not drawn on fig 5)). | As our instrument is designed for a linear input polarization, a polarizing cube is inserted in the doublet spacing in order to select and fixes a linear vertical input polarization (not drawn on fig \ref{schema_complet}) ). |
The experimental setup can be seen in fig.6.. | The experimental setup can be seen in \ref{photo_objet}. |
The telescope array arrangement has to be carefully selected. to fit the sampling criteria for a proper image analysis. | The telescope array arrangement has to be carefully selected to fit the sampling criteria for a proper image analysis. |
As previously demonstrated (Armancletal.2008).. high dynamics imaging capability requires a recluncant array configuration. | As previously demonstrated \citep{A}, high dynamics imaging capability requires a redundant array configuration. |
Consequently. our telescope array must periodically sample the spatial frequency domain. | Consequently, our telescope array must periodically sample the spatial frequency domain. |
The object dimension and the focal length of the collimator have to be determined by comparing the object spectrum and the spatial frequencies sampled by our instrument. | The object dimension and the focal length of the collimator have to be determined by comparing the object spectrum and the spatial frequencies sampled by our instrument. |
The intensity ανν) observed in the image plane of the instrument isgiven by : | The intensity $I(At;y)$ observed in the image plane of the instrument isgiven by : |
Of the 17 galaxies in the Noo et al. (1995)) | Of the 17 galaxies in the Koo et al. \cite{koo95}) ) |
sample of ‘dnt blue galaxies. niue have redshifts such that the CO J-2-l aud J=3-2 lines are accessible with the IRAM receivers. | sample of faint blue galaxies, nine have redshifts such that the CO J=2-1 and J=3-2 lines are accessible with the IRAM receivers. |
Three of these nine galaxies are somewhat nore extended than the other six aud iudeed appear ion-stellar iu eround-based. optical observations (lXoo et al. 1995)). | Three of these nine galaxies are somewhat more extended than the other six and indeed appear non-stellar in ground-based optical observations (Koo et al. \cite{koo95}) ). |
We excluded these three. galaxies frou our siuuple on the erounds that they may be somewhat more nassive objects; akin perhaps to small spiral ealaxics. | We excluded these three galaxies from our sample on the grounds that they may be somewhat more massive objects, akin perhaps to small spiral galaxies. |
Observations of five of the remaining six galaxies were obtained with the IRAM 30 12 telescope in two separate observing runs in January 1996 aud June 1997. | Observations of five of the remaining six galaxies were obtained with the IRAM 30 m telescope in two separate observing runs in January 1996 and June 1997. |
The sixth galaxy. SÀ57-5182. was not observed due to tine constraints, | The sixth galaxy, SA57-5482, was not observed due to time constraints. |
The halfpower beam width is 17" at 2uni aud 12" at 1.3nuu. | The half-power beam width is $^{\prime\prime}$ at 2--mm and $^{\prime\prime}$ at 1.3–mm. |
We used the 2 and 1.3nuu SiS receivers to observe both lines simultaneously. | We used the 2– and 1.3–mm SiS receivers to observe both lines simultaneously. |
The receivers were all used iu single sideband mode. aud the vpieal svstein temperatures were 250-500 [KK for the 2uni receiver and 350-700 Is. for the 1.9nuu receiver. in M scale (or. ou average.S GOO IS ar 1200 I in Tap scale. respectively). | The receivers were all used in single sideband mode, and the typical system temperatures were 250-500 K for the 2--mm receiver and 350-700 K for the 1.3–mm receiver, in $_A^*$ scale (or, on average, 600 K and 1200 K in $_{MB}$ scale, respectively). |
The backeuds were esseutiallv two 1MIIZ-filter-bauks. of 512 channels cach. and in addition au auto-correlator: the spectra have been snoothed to 10 km 1 resolution. | The backends were essentially two 1MHz-filter-banks, of 512 channels each, and in addition an auto-correlator; the spectra have been smoothed to 10 km $^{-1}$ resolution. |
The observations were made using a nutatiug secondary with a beam throw of 1.5’. The poiuting was checked every two hours aud the Εμ accuracy was estimated to be 3” rns. | The observations were made using a nutating secondary with a beam throw of $^\prime$ The pointing was checked every two hours and the pointing accuracy was estimated to be $^{\prime\prime}$ rms. |
The spectra were first iuspected. aud any spectruii showing baseline. curvature or other artifacts was discarded. | The spectra were first inspected, and any spectrum showing baseline curvature or other artifacts was discarded. |
The remaining spectra were averaged togetlier. weighted by their rus noise. | The remaining spectra were averaged together, weighted by their rms noise. |
A first order bascline was removed from cach average spectrin. and the spectra were smoothed to a resolution of 10 kins |! to produce the final spectra (Fie. 1)). | A first order baseline was removed from each average spectrum, and the spectra were smoothed to a resolution of 10 km $^{-1}$ to produce the final spectra (Fig. \ref{fig-1}) ). |
The final temperatures (aud rnis noise in Table 1) have heen converted to the Ta;p teirperature scale Gjape=0.15 at 230 Cz. 0.59 at 150 CGIIz). | The final temperatures (and rms noise in Table 1) have been converted to the $T_{MB}$ temperature scale $\eta_{MB} = 0.45$ at 230 GHz, 0.59 at 150 GHz). |
Upper liuits to the integrated CO intensity were derived using the riis noise measured from the CO spectra aud the velocity widths obtained frou measurements of optical ciission lines (soo et al. 1995)). | Upper limits to the integrated CO intensity were derived using the rms noise measured from the CO spectra and the velocity widths obtained from measurements of optical emission lines (Koo et al. \cite{koo95}) ). |
We adopt as the 3c upper limit to the CO intensity (Wiklind Combes 199bj). where o is the τις nolse in K measured in our 10 kin | channels AV is the velocity width of the CO line. here taken to be the balfanaxinuun of the optical lines. and IN44 lans Lis the umuber of chaunels in the velocity width. | We adopt as the $\sigma$ upper limit to the CO intensity (Wiklind Combes \cite{wik94b}) ), where $\sigma$ is the rms noise in K measured in our 10 km $^{-1}$ channels, $\Delta V$ is the velocity width of the CO line, here taken to be the full-width half-maximum of the optical lines, and $N_{chan} = \Delta V/10$ km $^{-1}$ is the number of channels in the velocity width. |
The CO luminosity for a source at hieli redshift is eiven by where 5 is the area of the main beam iu square areseconds aud Dp=(eTT?\dotEUMLTT24021)| is the huninosity distance in Mpc (Wikliud Combes 1991b)). | The CO luminosity for a source at high redshift is given by where $\Omega_B$ is the area of the main beam in square arcseconds and $D_L = (c/H_o q_o^2)[q_o z + (q_o-1)(\sqrt{1+2q_oz}-1)]$ is the luminosity distance in Mpc (Wiklind Combes \cite{wik94b}) ). |
We adopt q,=0.5 aud 7L,=70 kin E E in this paper. | We adopt $q_o=0.5$ and $H_o = 70$ km $^{-1}$ $^{-1}$ in this paper. |
Table 1 gives the position. redshift. aud velocity width obtained from the optical cussion lines (Isoo ct al. 1995)). | Table \ref{tbl-1} gives the position, redshift, and velocity width obtained from the optical emission lines (Koo et al. \cite{koo95}) ), |
as well as the iuteeration time. the nus noise for cach line. aud the CO integrated intensity and CO luminosity calculated from the CO J=3-2 upper limit. | as well as the integration time, the rms noise for each line, and the CO integrated intensity and CO luminosity calculated from the CO J=3-2 upper limit. |
Our upper limits to the CO flux are comparable to the best upper limits in the literature for moderate to high redshift objects. | Our upper limits to the CO flux are comparable to the best upper limits in the literature for moderate to high redshift objects. |
For example. the detections of CO J=3-2 emission at hieh redshift are 6.7 Jv hans 1 for IRAS F1021111721 (Radford ct al. 1996)) | For example, the detections of CO J=3-2 emission at high redshift are 6.7 Jy km $^{-1}$ for IRAS F10214+4724 (Radford et al. \cite{radford}) ) |
and 8.1 Jy laus. ! for the Cloverleaf (Barvainis ct al. 1991)). | and 8.1 Jy km $^{-1}$ for the Cloverleaf (Barvainis et al. \cite{barvainis}) ), |
while our 30 upper linits rauge from 2 to 8 Jv lau +. | while our $\sigma$ upper limits range from 2 to 8 Jy km $^{-1}$. |
If ow ealaxies had comparable CO fluxes to IRAS F102111172 lor the Cloverleaf quasar. we would have detected them with our observations. | If our galaxies had comparable CO fluxes to IRAS F10214+4724 or the Cloverleaf quasar, we would have detected them with our observations. |
In addition. if we asstme the amplification duc to lensing is a factor of 10 in the two hieh redshift ealaxies. their CO huninosities Lew (converted to our cosinologyv) are 8.9<10? and 1.31019 Wy lan »pe. respectirverv. | In addition, if we assume the amplification due to lensing is a factor of 10 in the two high redshift galaxies, their CO luminosities $L_{CO}$ (converted to our cosmology) are $8.9\times 10^9$ and $1.3\times 10^{10}$ K km $^{-1}$ $^2$, respectively. |
. Thus.; we would have detected either of these wo galaxies. uuleused. at a redshift of 2~0.5. | Thus, we would have detected either of these two galaxies, unlensed, at a redshift of $z \sim 0.5$. |
Since these faint blue galaxies are thought o be distaut counterparts o UTM galaxies. we should also compare our upper lanits with CO observations of nearby ον ealaxies. | Since these faint blue galaxies are thought to be distant counterparts to HII galaxies, we should also compare our upper limits with CO observations of nearby dwarf galaxies. |
The CO J=1-0 luminosities of the starburst ealaxy M82 aud the IIII galaxy. UME18 are both Lew—510 Nus ! pe? (ealeulated from Young et al. 1995: | The CO J=1-0 luminosities of the starburst galaxy M82 and the HII galaxy UM448 are both $L_{CO} \sim 5\times 10^8$ K km $^{-1}$ $^2$ (calculated from Young et al. \cite{young}; |
Sage et al. 1992)). | Sage et al. \cite{sage}) ). |
Unfortunately. our ost upper limits are still a factor of LS larecr than the buuinosities of these nearby dwarf galaxies. aud so we would not have detected AIS2 or UALIS at 2~0.5. | Unfortunately, our best upper limits are still a factor of 4-8 larger than the luminosities of these nearby dwarf galaxies, and so we would not have detected M82 or UM448 at $z\sim 0.5$. |
For galaxies iu the local universe with ucar-solar metallicities and normal rates of star formation (1.6.not starburst galaxies). the mass of molecular hydrogen gas is related to the CO luminosity iu the J=1-0 line by Adj,=LSLeo Gc. Solomon et al. 1987)). | For galaxies in the local universe with near-solar metallicities and normal rates of star formation (i.e.not starburst galaxies), the mass of molecular hydrogen gas is related to the CO luminosity in the J=1-0 line by $M_{H_2} = 4.8 L_{CO}$ $_\odot$ (i.e. Solomon et al. \cite{sol87}) ). |
Siuce we have observed the CO J—2-1 aud J=3-2 lines. we lust consider the excitation of the eas in estimating nolecular eas masses. | Since we have observed the CO J=2-1 and J=3-2 lines, we must consider the excitation of the gas in estimating molecular gas masses. |
In ealactic nuclei. the three trausitious have | In galactic nuclei, the three transitions have |
is Comparable with the extent of the virialized (relaxed) core of a cluster. | is comparable with the extent of the virialized (relaxed) core of a cluster. |
Analyses of x-ray observations and dvnamical calculations on these scales make a variety of equilibrium and svinmeltry assumptions. | Analyses of x-ray observations and dynamical calculations on these scales make a variety of equilibrium and symmetry assumptions. |
Generally. agreement among (he various mass estimation techniques on this scale is impressive. | Generally, agreement among the various mass estimation techniques on this scale is impressive. |
Although there are still puzzles about clusters and their evolution. their central regions are reasonably well-studied over a wide redshift range. | Although there are still puzzles about clusters and their evolution, their central regions are reasonably well-studied over a wide redshift range. |
Many fewer observational studies have addressed the infall region thiat marks the transition between the cluster core and the surrounding large-scale structure. | Many fewer observational studies have addressed the infall region that marks the transition between the cluster core and the surrounding large-scale structure. |
At least in part. this inattention reflects the observational challenges of observing these larger. less dense regions. | At least in part, this inattention reflects the observational challenges of observing these larger, less dense regions. |
Now wilh wide-field spectroscopic instruments like the Ilectospec on the MALT (Fabricant et al. | Now with wide-field spectroscopic instruments like the Hectospec on the MMT (Fabricant et al. |
1998: Fabricant et al. | 1998; Fabricant et al. |
2005). it is possible to acquire dense samples of these fascinating regions (hat lie between Iso, and Ry... the radius of the shell of material just turning around from the ILIubble flow at redshift z (Gunn Gott 1972: Kaiser 1987; Reeos Geller 1939). | 2005), it is possible to acquire dense samples of these fascinating regions that lie between $_{200}$ and $_{turn}$, the radius of the shell of material just turning around from the Hubble flow at redshift $z$ (Gunn Gott 1972; Kaiser 1987; Regos Geller 1989). |
The infall region is a route to understanding the growth rate of clusters. their ultimate masses. and the relationship between galaxy and cluster evolution (Dialerio Geller 1997: Ellingson et al. | The infall region is a route to understanding the growth rate of clusters, their ultimate masses, and the relationship between galaxy and cluster evolution (Diaferio Geller 1997; Ellingson et al. |
2001: Busha et al. | 2001; Busha et al. |
2005: Rines et al. | 2005; Rines et al. |
2005: Tran et al. | 2005; Tran et al. |
2005). | 2005). |
On the scale of the infall region. there are only two techniques to probe the matter distribution. weak lensing (e.g. Lemze et al. | On the scale of the infall region, there are only two techniques to probe the matter distribution, weak lensing (e.g. Lemze et al. |
2009: Umetsu et al. | 2009; Umetsu et al. |
2011) and a kinematic technique called the caustic method (Diaferio Geller 1997: Dialerio 1999: Serra et al. | 2011) and a kinematic technique called the caustic method (Diaferio Geller 1997; Diaferio 1999; Serra et al. |
2011). | 2011). |
Neither of these methods depends on the dynamical state of the svstem and both apply at all clustrocentric radii (Diaferio. Geller Rines 2005). | Neither of these methods depends on the dynamical state of the system and both apply at all clustrocentric radii (Diaferio, Geller Rines 2005). |
Ol course. for nearly all clusters. we can observe them only in redshift (phase) space. | Of course, for nearly all clusters, we can observe them only in redshift (phase) space. |
Kaiser (1987) was the first to understand. how spherical infall appears in redshilt space. | Kaiser (1987) was the first to understand how spherical infall appears in redshift space. |
In his elegant paper (Ixaiser 1937). he shows (his Figure 5) the now widely recognized pattern that characterizes (he appearance of a cluster in redshift space. | In his elegant paper (Kaiser 1987), he shows (his Figure 5) the now widely recognized trumpet-shaped pattern that characterizes the appearance of a cluster in redshift space. |
The central. virialized region appears as an extended finger pointing along the lime-of-sight toward the observer. | The central, virialized region appears as an extended finger pointing along the line-of-sight toward the observer. |
This elongation is a simple consequence of the fact that the line-ol-sielt component ol the velocities of galaxies relative to one another within the virialized region are larger than the Lhabble flow across the region. | This elongation is a simple consequence of the fact that the line-of-sight component of the velocities of galaxies relative to one another within the virialized region are larger than the Hubble flow across the region. |
At the effective outer radius of the cluster. Ry... the infall velocily just cancels the IIubble flow. | At the effective outer radius of the cluster, $_{turn}$, the infall velocity just cancels the Hubble flow. |
Thus the shell just turning around appears as a line ab (he cluster mean velocity in redshilt space. | Thus the shell just turning around appears as a line at the cluster mean velocity in redshift space. |
Infalling shells at radii between δεν anc Rey are successively more and more elongated along the line-of-sight producing the trumpet shape. | Infalling shells at radii between $_{turn}$ and $_{200}$ are successively more and more elongated along the line-of-sight producing the trumpet shape. |
Ii the simple spherical infall model. the outline of the trumpet is a (rue caustic (a line of infinite density in phase space). | In the simple spherical infall model, the outline of the trumpet is a true caustic (a line of infinite density in phase space). |
At about the same time that Ixaiser wrote his paper. there was an increasing awareness | At about the same time that Kaiser wrote his paper, there was an increasing awareness |
his paper. | this paper. |
The fine-structure lines are. however. a valuable ool to infer the physical conditions in such svstenis. | The fine-structure lines are, however, a valuable tool to infer the physical conditions in such systems. |
In xwlicular. the knowledge of the ionization. state of. the sVslers couped with the information on the volumentric density alloreed by the fine-structure lines allows one to ace limits οn the distance between the absorber ancl the OSO. giving a clue to infer whether they correspond to intervening couds or to material ejected from the QSO (lurnshek.\Wevmann&Williams1979:Morrisetal.1986:Tripp.Lu&Savage1996:Srianand.Petitjean 2000). | In particular, the knowledge of the ionization state of the systems coupled with the information on the volumentric density afforded by the fine-structure lines allows one to place limits on the distance between the absorber and the QSO, giving a clue to infer whether they correspond to intervening clouds or to material ejected from the QSO \cite{TWW79,Morris,TLS96,SP2000}. |
. So lar. all the fine-structure lines observed. belong to cither C" or C... | So far, all the fine-structure lines observed belong to either $^0$ or $^+$. |
Owing to its low ionization fraction (since its lonization potential is lower than that of. hydrogen). atomic carbon is very seldom detected. | Owing to its low ionization fraction (since its ionization potential is lower than that of hydrogen), atomic carbon is very seldom detected. |
The three systems listed in table 2) correspond to all of the presently known svstems. apart from the system observed. towards the BL Lac object 0215|015 (Bladesetal.1982:Blades 1985). | The three systems listed in table \ref{obsdata} correspond to all of the presently known systems, apart from the system observed towards the BL Lac object 0215+015 \cite{Bladesa,Bladesb}. |
As we gathered observational data from the literature. we rejected any line falling within the Ly-a forest region of the spectrum. | As we gathered observational data from the literature, we rejected any line falling within the $\alpha$ forest region of the spectrum. |
Prochaska (1999) observed. the 1335 fine-structure transition in a LL system at sn.=2.652 towards Q2231-00. | Prochaska \shortcite{P99} observed the 1335 fine-structure transition in a LL system at $z_{\rmn{abs}}=2.652$ towards Q2231-00. |
Llowever. since this transition falls within the Ly-a forest in this object and therefore may have been subject to significant contamination. his claimed value on the column density N(C 1) should be regarded at most as an upper limit to the true value. | However, since this transition falls within the $\alpha$ forest in this object and therefore may have been subject to significant contamination, his claimed value on the column density $^*$ ) should be regarded at most as an upper limit to the true value. |
For the same reason we disregarded the DLA system at zi=3.054 towards QO000-26 observed by Giardino Favata (2000). | For the same reason we disregarded the DLA system at $z_{\rmn{abs}}=3.054$ towards Q0000-26 observed by Giardino Favata \shortcite{GF2000}. |
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