source
stringlengths 1
2.05k
⌀ | target
stringlengths 1
11.7k
|
|---|---|
The “cooling length scale" J, is given by an optical depth of t=1 in the long wavelength regime, and the “irradiation length scale" /, is given by an optical depth of 7=1 the broad stellar irradiation spectrum. | The “cooling length scale” $l_\mathrm{c}$ is given by an optical depth of $\tau = 1$ in the long wavelength regime, and the “irradiation length scale” $l_*$ is given by an optical depth of $\tau = 1$ the broad stellar irradiation spectrum. |
Fig. | Fig. |
1 shows the cooling length scale in the gray approximation as a function of the location of the cavity shell Reayity. | \ref{fig:ResolutionIR} shows the cooling length scale in the gray approximation as a function of the location of the cavity shell $R_\mathrm{cavity}$. |
The shell density is chosen to be p=4x107!"gcm” as depicted in Fig. | The shell density is chosen to be $\rho = 4 \times 10^{-17} \rhocgs$ as depicted in Fig. |
5. | 5. |
The opacities are computed as the Rosseland mean opacities of ?.. | The opacities are computed as the Rosseland mean opacities of \citet{Laor:1993p736}. |
The stellar temperature is taken from the ? tracks for a star and an accretion rate of M=10?Mo γι, and the temperature at the location Reayity is computed with a slope ΤocΑνν in the gray approximation. | The stellar temperature is taken from the \citet{Hosokawa:2009p12591} tracks for a star and an accretion rate of $\dot{M} = 10^{-3} \Msol \mbox{ yr}^{-1}$ , and the temperature at the location $R_\mathrm{cavity}$ is computed with a slope $T \propto R_\mathrm{cavity}^{-0.5}$ in the gray approximation. |
For simplification, the evaporation of dust grains is neglected, but wouldresult in an even larger cooling length-scale. | For simplification, the evaporation of dust grains is neglected, but wouldresult in an even larger cooling length-scale. |
Hence, we can be sure that the length scale of cooling is resolved. | Hence, we can be sure that the length scale of cooling is resolved. |
The absorption length scale /. the stellar irradiation spectrum is much smaller than the cooling length scale for thermal dust emission. | The absorption length scale $l_*$ the stellar irradiation spectrum is much smaller than the cooling length scale for thermal dust emission. |
Using again the opacities of ?,, Fig. | Using again the opacities of \citet{Laor:1993p736}, , Fig. |
2 shows the absorption length scale the whole stellar irradiation spectrum for different gas densities. | \ref{fig:Resolution} shows the absorption length scale the whole stellar irradiation spectrum for different gas densities. |
Cradieunts in Faraday rotation across turbulent cells can lead to correlations between rotativitv aud Stokes Q aud C parameters. which can result in “correlation depolarization”. | Gradients in Faraday rotation across turbulent cells can lead to correlations between rotativity and Stokes $Q$ and $U$ parameters, which can result in “correlation depolarization”. |
Observed polarization levels require that he feld have niu reversals along the Lue of sight to avoid this effect. | Observed polarization levels require that the field have many reversals along the line of sight to avoid this effect. |
Statistical fluctuations of circular aud incar polarizations are then likely to be dominated bv changes in the mean paramcters describing the plasma rather than by the stochastic behavior of the turbuleut uediunu. | Statistical fluctuations of circular and linear polarizations are then likely to be dominated by changes in the mean parameters describing the plasma rather than by the stochastic behavior of the turbulent medium. |
Variations iu the mean parameters are unlikely o change the helicity of circular polarization unless a source undergoes a sharp transition from very low to very Heh svuchrotron We have shown that our model is potentially applicable to a wide range of compact svuchrotron sources. | Variations in the mean parameters are unlikely to change the helicity of circular polarization unless a source undergoes a sharp transition from very low to very high synchrotron We have shown that our model is potentially applicable to a wide range of compact synchrotron sources. |
Iu particular. it naturallv predicts an excess of circular over linear polarization when a source is strongly depolarized. by the mean Faraday rotation and when a small απλο! of linear polarizarion is cficicutly converted mto circular polarization. | In particular, it naturally predicts an excess of circular over linear polarization when a source is strongly depolarized by the mean Faraday rotation and when a small amount of linear polarizarion is efficiently converted into circular polarization. |
This can explain the polarization properties of the Galactic Center and M8SI. | This can explain the polarization properties of the Galactic Center and $^{*}$. |
We thank Roger Blandford. Avery Broderick and Marck Sikora for insightful discussions. | We thank Roger Blandford, Avery Broderick and Marek Sikora for insightful discussions. |
This work was supported in part bv NSF eraut AST-08TGSST. | This work was supported in part by NSF grant AST-9876887. |
this work. which have [fixie] as laree as 200. are thus not in conllict with the CAIB on the scales which are relevant [or our analysis. | this work, which have $|f_{\rm NL}|$ as large as 200, are thus not in conflict with the CMB on the scales which are relevant for our analysis. |
A very promising wav to constrain departures [rom CGaussianity is to measure the various properties of massive virialized. structures like their abundance (2?77).. clustering and their biasing (?7??????7).. | A very promising way to constrain departures from Gaussianity is to measure the various properties of massive virialized structures like their abundance \citep{matarrese00, VJKM01, loverde08}, clustering and their biasing \citep{GW86, MLB86,
matarrese08,DDHS07,carbone08,seljak08,matarrese08}. |
Indeed. the best constraints on non-CGaussianitv [rom the LSS have been obtained. by 7 including the observed. scale-dependent bias of. the spectroscopic sample SDSS luminous rec galaxies and. the whotometric quasar sample. | Indeed, the best constraints on non-Gaussianity from the LSS have been obtained by \cite{slosar08} including the observed scale-dependent bias of the spectroscopic sample SDSS luminous red galaxies and the photometric quasar sample. |
“Phe resulting limits of 29<fe«10 (95 confidence level) are remarkably close to those obtained rom the CAIB analysis alone and. according to ?.. could oe further improved by looking for scale dependency in he biasing of two different population of objects. | The resulting limits of $-29 < f_{\rm NL} < 70$ (95 confidence level) are remarkably close to those obtained from the CMB analysis alone and, according to \cite{seljak08}, could be further improved by looking for scale dependency in the biasing of two different population of objects. |
Alternatively, one can consider the topology of the mass density Ποια (2).. and higher-order clustering statistics like he bispeetrum (2).. | Alternatively, one can consider the topology of the mass density field \citep{matsubara03}, , and higher-order clustering statistics like the bispectrum \citep{hikage06}. |
Phe ability of these techniques to detect he imprint. of the primordial non-Gaussianity on the LSs ias been tested with N-body experiments (272??22227). | The ability of these techniques to detect the imprint of the primordial non-Gaussianity on the LSS has been tested with N-body experiments \citep{messina90,moscardini91,weinbergcole92,mathis04,kang07,grossi07,DDHS07,hikage08}. |
N-»ody simulations are of paramount importance in the study of NC models. since one needs to ciscntangle primordial non-Gaussianity from late non-CGaussianity. induced by the non-linear growth of density perturbations that can only be properly accounted for by numerical experiments. | N-body simulations are of paramount importance in the study of NG models, since one needs to disentangle primordial non-Gaussianity from late non-Gaussianity induced by the non-linear growth of density perturbations that can only be properly accounted for by numerical experiments. |
Reeenthy. 2? have carried. out. cosmological N-bocky simulations of NC mocels to study the evolution of the probability distribution. function. (PDE) of the density Iuctuations. | Recently, \cite{grossi07,grossi08} have carried out cosmological N-body simulations of NG models to study the evolution of the probability distribution function (PDF) of the density fluctuations. |
Phey found that the imprint of primordial non-CGaussianity. which is evident in the negative tail of the PDE at high redshifts. is preserved throughout the subsequent evolution and out to the present epoch. | They found that the imprint of primordial non-Gaussianity, which is evident in the negative tail of the PDF at high redshifts, is preserved throughout the subsequent evolution and out to the present epoch. |
Εις result suggests that void statistics may be a promising effective tool for detecting primordial non-Caussianity (77)) and that can be applied to dillerent tvpes of observations over a large range of cosmic epochs. | This result suggests that void statistics may be a promising effective tool for detecting primordial non-Gaussianity \cite{kamionkowski08,song08}) ) and that can be applied to different types of observations over a large range of cosmic epochs. |
Taking advantage of the recent theoretical cllorts for standardizing the appropriate statistical tools (2) one could apply void-finding algorithms to quantify the properties of the unclerdense regions observed in the spatial distribution of galaxies. | Taking advantage of the recent theoretical efforts for standardizing the appropriate statistical tools \citep{colberg08} one could apply void-finding algorithms to quantify the properties of the underdense regions observed in the spatial distribution of galaxies. |
Unfortunately. current galaxy recdshift surveys are. probably. too small for void-based statistics to appreciate deviations from the Gaussian case at the level required. | Unfortunately, current galaxy redshift surveys are probably too small for void-based statistics to appreciate deviations from the Gaussian case at the level required. |
“Phe situation will change in a not too distant future. when next generation all-sky surveys like ADEPT or EUCLID will allow to measure he position of ~5«10' ogalaxies over a largeὃν rangeo of redshifts out to 2= 2. | The situation will change in a not too distant future, when next generation all-sky surveys like ADEPT or EUCLID will allow to measure the position of $\sim 5\times 10^7$ galaxies over a large range of redshifts out to $z=2$ . |
Alternatively. once can analyze igh-resolution spectra of distant quasars to characterize he properties of the underlving mass censity field. at Lol3 (eg. ??7)). | Alternatively, one can analyze high-resolution spectra of distant quasars to characterize the properties of the underlying mass density field at $z>3$ (e.g. \cite{viel03,viel04,vielbispect,lesg}) ). |
In particular. since we expect that underdense regions are characterized by a low neutral ivdrogen (LL) abundance. one can infer the presence of voids and quantify their statistical properties from. voids in the ransmitted flux. defined as the connected. regions in the spectral Lux distribution above the mean Εαν level. | In particular, since we expect that underdense regions are characterized by a low neutral hydrogen (HI) abundance, one can infer the presence of voids and quantify their statistical properties from voids in the transmitted flux, defined as the connected regions in the spectral flux distribution above the mean flux level. |
The connection between voids and spectral regions characterized » negligible HE absorption has been recently studied by C3 using hvdrodynamical simulations where a link at z~ oetween the [ux and matter properties is provided. | The connection between voids and spectral regions characterized by negligible HI absorption has been recently studied by \cite{viel08voids} using hydrodynamical simulations where a link at $z\sim 2$ between the flux and matter properties is provided. |
In this work we perform. for the [first time. high-resolution hyerocsynamical. simulations of NC models. to heck whether one can use the intergalactic medium (see ? for a recent review) to detect non-Gaussian features in 16 [lux statistics like the PDE. ux power and 10 bispectrum. | In this work we perform, for the first time, high-resolution hydrodynamical simulations of NG models to check whether one can use the intergalactic medium (see \cite{meiksin07} for a recent review) to detect non-Gaussian features in the flux statistics like the PDF, flux power and the bispectrum. |
The lavout of the paper is as follows. | The layout of the paper is as follows. |
In Vlection 2 we describe the hvdrodynamical simulations and we show an example of simulated quasar (QSO) μα»ecteuni. | In Section 2 we describe the hydrodynamical simulations and we show an example of simulated quasar (QSO) spectrum. |
In Section 3 we present the results of the various —ux statistics. | In Section 3 we present the results of the various flux statistics. |
In Section 4 we address the role of svstematic ancl statistical errors that could contaminate the NC signal. | In Section 4 we address the role of systematic and statistical errors that could contaminate the NG signal. |
We conclude in Section 5. | We conclude in Section 5. |
We rely on simulations run with the parallel hvdrodynamical CIreeSPLHI) codeGADGLEL-2 based on the conservative "entropy-formulation of SPIL (7).. | We rely on simulations run with the parallel hydrodynamical (TreeSPH) code based on the conservative `entropy-formulation' of SPH \citep{springel}. |
They consist. of a cosmological volume with periodic boundary conditions filled with an equal number of dark matter and gas particles. | They consist of a cosmological volume with periodic boundary conditions filled with an equal number of dark matter and gas particles. |
ltadiative cooling and heating processes were followed: for a primordial mix of hydrogen and. helium. | Radiative cooling and heating processes were followed for a primordial mix of hydrogen and helium. |
We assumed a mean Ultraviolet Backerouncl similar to that propesed by ? produced by quasars and. galaxies as given by with helium. heating rates multiplied by a factor 3.3 in order to better fit observational constraints on the temperature evolution of the ICM (e.g. 22)). | We assumed a mean Ultraviolet Background similar to that propesed by \cite{haardt1996} produced by quasars and galaxies as given by with helium heating rates multiplied by a factor 3.3 in order to better fit observational constraints on the temperature evolution of the IGM (e.g. \cite{schaye00,ricotti00}) ). |
This background. gives naturally a hvdrogen ionization rate Poy.~1 at the redshifts of interest. here (e.g. 22)). | This background gives naturally a hydrogen ionization rate $\Gamma_{-12}\sim 1$ at the redshifts of interest here (e.g. \cite{bolt05,fg08}) ). |
The star formation criterion is a very simple one that converts in collisionless stars all the gas particles whose temperature falls below 107 Ix and whose density contrast is larger than 1000 (it has been shown that the star formation criterion has a negligible impact on Dux statistics). | The star formation criterion is a very simple one that converts in collisionless stars all the gas particles whose temperature falls below $10^5$ K and whose density contrast is larger than 1000 (it has been shown that the star formation criterion has a negligible impact on flux statistics). |
More details can be found in (?).. | More details can be found in \citep{viel04}. |
The cosmological reference. model. corresponds. to à ‘fiducial [CDM Universe with parameters. at z=0. Ou=0.26.OX=0.74.Οι, 0.0463. n,=0.95. and 44,=72 km + Lande=0.85 (the B2 series of 2)). | The cosmological reference model corresponds to a `fiducial' $\Lambda$ CDM Universe with parameters, at $z=0$, $\Omega_{\rm m
}=0.26,\ \Omega_{\rm \Lambda}=0.74,\ \Omega_{\rm b }=0.0463$ , $n_{\rm
s}=0.95$, and $H_0 = 72$ km $^{-1}$ $^{-1}$ and $\sigma_8=0.85$ (the B2 series of \cite{viel04}) ). |
We have used 2.384? dark matter and gas particles in a 60fh+ comoving Alpe box for the flux power ancl bispectrum. to better sample the large scales. | We have used $2\times 384^3$ dark matter and gas particles in a $60\ h^{-1}$ comoving Mpc box for the flux power and bispectrum, to better sample the large scales. |
For the flux probability clistribution function we relied instead on 2«256" dark matter and gas particles. ina ..20.17 comoving: Alpe. since. below and around =3 this seems to be the appropriate resolution the ect numerical convergence. | For the flux probability distribution function we relied instead on $2\times
256^3$ dark matter and gas particles in a $20\ h^{-1}$ comoving Mpc, since below and around $z=3$ this seems to be the appropriate resolution the get numerical convergence. |
The gravitational softening was set to 2.5 and 5 f kpe in comoving units for all particles for the 20 and 60 comoving Alpesh boxes. respectively. | The gravitational softening was set to 2.5 and 5 $h^{-1}$ kpc in comoving units for all particles for the 20 and 60 comoving $h$ boxes, respectively. |
The mass per gas particle is 6.12«10" M.fh for the small boxes and 4.9« 10M./h [or the Large boxes. while the high resolution run for the small box has à mass per gas particle of LS.10 ML.ff (this refers to a (20.384) simulation that was performed in order to check for numerical convergence of the Dux PDE). | The mass per gas particle is $6.12\times10^6 $ $_{\odot}/h$ for the small boxes and $4.9\times10^7 $ $_{\odot}/h$ for the large boxes, while the high resolution run for the small box has a mass per gas particle of $1.8\times10^6 $ $_{\odot}/h$ (this refers to a (20,384) simulation that was performed in order to check for numerical convergence of the flux PDF). |
In the following. the different simulations will be indicated by two numbers. (Np.We): Ny is the size of the box in comoving Alpe/f and No is the cubie root of the totalnumber of gas particles in the simulation. | In the following, the different simulations will be indicated by two numbers, $(N_1,N_2)$: $N_1$ is the size of the box in comoving $/h$ and $N_2$ is the cubic root of the totalnumber of gas particles in the simulation. |
NG are produced in the initial conditions at z=99 using the same method as in ? that we brielly summarize here. | NG are produced in the initial conditions at $z=99$ using the same method as in \cite{grossi07} that we briefly summarize here. |
Initial NC conditions are generatecl without mocdifvingthe linear matter power spectrum using the Zel'dovieh approximation: a Ciaussian gravitational potential is generated in Fourier space from a power-law power spectrum of the form P(A)x&7 and | Initial NG conditions are generated without modifyingthe linear matter power spectrum using the Zel'dovich approximation: a Gaussian gravitational potential is generated in Fourier space from a power-law power spectrum of the form $P(k)\propto
k^{-3}$ and |
quoted above (Grillin et al. | quoted above (Griffin et al., |
1986: Orton et al. | 1986; Orton et al., |
1986: Crillin Orton 1993) | 1986; Griffin Orton 1993). |
Based on this. we expect the peak [lux density of Mars to be 8.1 Jv/beam. | Based on this, we expect the peak flux density of Mars to be 8.1 kJy/beam. |
During the observations the measured value of merc. was 0.0455. | During the observations the measured value of $\tau_{225 GHz}$ was 0.0455. |
Using our relation above. between 1.5 Pilz and 225 6111. this corresponds to à value of τισ of 4.32. | Using our relation above, between 1.5 THz and 225 GHz, this corresponds to a value of $\tau_{1.5 THz}$ of $\sim$ 4.32. |
Therefore the observed Dux density at the telescope is estimated to be 30 Jv/beam. | Therefore the observed flux density at the telescope is estimated to be 30 Jy/beam. |
The integration time per point of the map was 50 seconds. | The integration time per point of the map was 50 seconds. |
The peak was detected at a level of 73.2 o. | The peak was detected at a level of $\sim$ 3.2 $\sigma$. |
Fherefore we calculate from the Mars data that the noise equivalent (ux density (NERD) of the TIIUMPELR. combination is ~66410 Jy (10 Is). | Therefore we calculate from the Mars data that the noise equivalent flux density (NEFD) of the JCMT-THUMPER combination is $\sim$ $\pm$ 10 Jy $\sigma$ 1s). |
We note once again that our detection is not at the 70 level. | We note once again that our detection is not at the $\sigma$ level. |
However. we believe it is also a real detection for similar reasons to those quoted above: the source was seen in several pixels simultaneously: and the NEPD we calculate rom the measurements is consistent with that. predicted rom laboratory measurements of the detector svstem and with that seen in the Jupiter data. | However, we believe it is also a real detection for similar reasons to those quoted above: the source was seen in several pixels simultaneously; and the NEFD we calculate from the measurements is consistent with that predicted from laboratory measurements of the detector system and with that seen in the Jupiter data. |
The map of Mars shows a slightly dillerent morphology rom that of Jupiter. | The map of Mars shows a slightly different morphology from that of Jupiter. |
The map shows some evidence that he brightest region is more centrally peaked. and that this central peak sits on an extended plateau. | The map shows some evidence that the brightest region is more centrally peaked, and that this central peak sits on an extended plateau. |
Converselv. Jupiter is better fitted by a single gaussian. | Conversely, Jupiter is better fitted by a single gaussian. |
For Mars the centrally veaked core is of order 215 arcsec across. with a more extended lobe to the north. | For Mars the centrally peaked core is of order $\sim$ 15 arcsec across, with a more extended lobe to the north. |
Phe core is consistent with the FWHIAL that. would be expected. from. Mars. since at. the ime of the observation the diameter o£ Mars was 6 aresec. | The core is consistent with the FWHM that would be expected from Mars, since at the time of the observation the diameter of Mars was 6 arcsec. |
When convolved with our beam this produces a gaussian of JIWIIM 15.2 aresec. consistent with the image in Figure δ.. | When convolved with our beam this produces a gaussian of FWHM 15.2 arcsec, consistent with the image in Figure \ref{mars}. |
The extended emission to the north is most likely to be the οοσα of the telescope. | The extended emission to the north is most likely to be the error-beam of the telescope. |
The level of the extended error-Iobe is roughly two-thirds ofthe magnitude of the main beam. | The level of the extended error-lobe is roughly two-thirds of the magnitude of the main beam. |
This may also extend to the south and south-cast. but our map does not extend far enough in these directions to sav. | This may also extend to the south and south-east, but our map does not extend far enough in these directions to say. |
Once again the instrument focus may be adding to the problem. | Once again the instrument focus may be adding to the problem. |
By the time we repeated the map the shift was coming to an end and the sun had risen. | By the time we repeated the map the shift was coming to an end and the sun had risen. |
Consequently. the conditions worsened (the sky noise level increased) anc we did not detect. Mars again in day-time. | Consequently, the conditions worsened (the sky noise level increased) and we did not detect Mars again in day-time. |
Our mean NEED measured on the planets is therefore ~65+10 Jv (la Is). | Our mean NEFD measured on the planets is therefore $\sim$ $\pm$ 10 Jy $\sigma$ 1s). |
Once the Sun hack risen we pointed. the telescope at the Sun. when it was in the airmass range 1.ΕΕ to 1.43. and detected it clearly in all channels. at high levels of to-noise ratio of up to 560r o. | Once the Sun had risen we pointed the telescope at the Sun, when it was in the airmass range 1.44 to 1.43, and detected it clearly in all channels, at high levels of signal-to-noise ratio of up to 560 $\sigma$. |
Obviously. we did not map the full extent of the Sun. but rather used it as a bright. uniform. extended: black-body source. ancl chopped across the limb of the Sun. | Obviously, we did not map the full extent of the Sun, but rather used it as a bright, uniform, extended black-body source, and chopped across the limb of the Sun. |
We estimate that the emission. [rom the Sun is 1.1 Mvbeam. | We estimate that the emission from the Sun is $\sim$ 1.1 MJy/beam. |
I was at 71.4 airmasses during our observation. and 75064: was 0.047. | It was at $\sim$ 1.4 airmasses during our observation, and $\tau_{225GHz}$ was 0.047. |
We used 5-second integrations. so we estimate our NEED on the Sun to be Jv (lo Is). | We used 5-second integrations, so we estimate our NEFD on the Sun to be $\sim$ 9 Jy $\sigma$ 1s). |
‘This is a factor of ~7 better than our NEED estimated rom the essentially point-like planet. Mars. | This is a factor of $\sim$ 7 better than our NEFD estimated from the essentially point-like planet Mars. |
Ες implies hat only roughly one-seventh (15543) of the total power incident upon the ΟΝΕ dish at 200 sam is focussed. into a ld-arcsec central beam. | This implies that only roughly one-seventh $\sim$ ) of the total power incident upon the JCMT dish at 200 $\mu$ m is focussed into a 14-arcsec central beam. |
We had. predicted. that. due to he surface inaccuracies of the JCAL dish. Gvhich was not designed to operate at this high a frequency). there would »' significant. power in the sicle-lobes. | We had predicted that, due to the surface inaccuracies of the JCMT dish (which was not designed to operate at this high a frequency), there would be significant power in the side-lobes. |
The measurement of the surface accuracy of the dish which is closest in time to our Observations was taken on 2005 February 13. at which time the dish surface accuracy was found to be 23.8 pam (Wouterloot 2005). | The measurement of the surface accuracy of the dish which is closest in time to our observations was taken on 2005 February 13, at which time the dish surface accuracy was found to be 23.8 $\mu$ m (Wouterloot 2005). |
Using the standard: calculation of Ruze ellicieney we estimate that a dish. of this accuracy would. concentrate ~LI% of the total power into a I1H-aresec. beam. | Using the standard calculation of Ruze efficiency we estimate that a dish of this accuracy would concentrate $\sim$ of the total power into a 14-arcsec beam. |
This is consistent. with the ~15% we estimate here based on the ratio of NEEDs measurecl on Mars ancl the Sun. | This is consistent with the $\sim$ we estimate here based on the ratio of NEFDs measured on Mars and the Sun. |
Hence we see that all of our observations are selí-consistent. | Hence we see that all of our observations are self-consistent. |
We have successfully commissioned the TIIUMPERB camera ab JCAPPE. | We have successfully commissioned the THUMPER camera at JCMT. |
We have demonstrated that -μ astronomy 1s possible from the ground. | We have demonstrated that $\mu$ m astronomy is possible from the ground. |
We have taken the first. grounel-based. images of Jupiter and. Mars. at this wavelength. | We have taken the first ground-based images of Jupiter and Mars at this wavelength. |
We have calibrated the 200-/m sky. ancl found. a relation between opacities of Tjτην=(05+10)©7o»s6gz. | We have calibrated the $\mu$ m sky, and found a relation between opacities of $\tau_{1.5{\rm THz}} = (95 \pm 10) \times \tau_{225{\rm GHz}}$. |
This is consistent with previous measurements. and also with our modelling of atmospheric transmission across this section of the clectro-magnetic spectrum. | This is consistent with previous measurements, and also with our modelling of atmospheric transmission across this section of the electro-magnetic spectrum. |
We estimate the NEED of the JCNEE.PILUAPER. svstem is ~65 Jv (la Is). | We estimate the NEFD of the JCMT–THUMPER system is $\sim$ 65 Jy $\sigma$ 1s). |
Subsets and Splits
No saved queries yet
Save your SQL queries to embed, download, and access them later. Queries will appear here once saved.