ReadingTimeMachine/rtm-sgt-ocr-v1 · Datasets at Hugging Face

source stringlengths 1 2.05k ⌀	target stringlengths 1 11.7k
]t is important to have in münd that the process of structure. formation of the Universe. ds. a very complicated one.	It is important to have in mind that the process of structure formation of the Universe is a very complicated one.
For the canonical model of cosmology. namely primordial Gaussian random perturbation field. the formation of voids and clusters of galaxies is part of one process of clustering.	For the canonical model of cosmology, namely primordial Gaussian random perturbation field, the formation of voids and clusters of galaxies is part of one process of clustering.
The formation of regions with positive density contrasts. like clusters of galaxies. contributes to the erowth of underdense regions and. vice versa.	The formation of regions with positive density contrasts, like clusters of galaxies, contributes to the growth of underdense regions and vice versa.
The present study only. considers a small part. of the very complicated: structure formation. process.	The present study only considers a small part of the very complicated structure formation process.
We study the evolution of negative density perturbations taking into account all relevant physical processes acting on them.	We study the evolution of negative density perturbations taking into account all relevant physical processes acting on them.
Llowever. the assessment of the role of these very physical processes not considered in the literature. besides including the presence of non-barvonic dark matter ancl cosmological constant. are the maincontributions of the present studs.	However, the assessment of the role of these very physical processes not considered in the literature, besides including the presence of non-baryonic dark matter and cosmological constant, are the maincontributions of the present study.
idoes per square degree with aperture mass greater than Aly=0.04 and AL,=0.08 for the filter scale @=2’ or the five cosmological models. using two different redshift intervals. zO.15.04]) and. z0.4.1].	haloes per square degree with aperture mass greater than $M_{\rm ap}=0.04$ and $M_{\rm ap}=0.08$ for the filter scale $\theta=2'$ for the five cosmological models, using two different redshift intervals, $z\in[0.15,0.4]$ and $z\in[0.4,1]$.
By. comparing he halo densities in the different recshift intervals for the various cosmologies in refndachx. and 3... we expect the Iargest dillerences between he cosmological models for AZ,=0.08.	By comparing the halo densities in the different redshift intervals for the various cosmologies in \\ref{ndachx} and \ref{ndach1}, we expect the largest differences between the cosmological models for $M_{\rm ap}=0.08$.
The reason for this is the stronger evolution for the rich cluster mass function which corresponds to large values of the aperture mass (sce Figure 1)).	The reason for this is the stronger evolution for the rich cluster mass function which corresponds to large values of the aperture mass (see Figure \ref{map}) ).
Whereas the EdS(0.6.0.25) and EcdS(1.0.25) models are again very ilferent. from the other three. the use of redshift information greatly helps to distinguish the EdS(0.6.0.5) model from the two low-clensity niocels.	Whereas the EdS(0.6,0.25) and EdS(1,0.25) models are again very different from the other three, the use of redshift information greatly helps to distinguish the EdS(0.6,0.5) model from the two low-density models.
For the latter. a survey area of less than 3 deg? would be sullicient.	For the latter, a survey area of less than 3 ${\rm deg}^2$ would be sufficient.
One might think of another way to obtain reclshilt information. namely to use source galaxies at. cillerent redshilts (clistinguishecd. say. by photometric redshift estimates).	One might think of another way to obtain redshift information, namely to use source galaxies at different redshifts (distinguished, say, by photometric redshift estimates).
To investigate this effect. we have plottec in Figure 7. the dependence of the number of haloes on the redshift of the sources for AL,=0.04 and 8=2’.	To investigate this effect, we have plotted in Figure \ref{sredshift} the dependence of the number of haloes on the redshift of the sources for $M_{\rm ap}=0.04$ and $\theta=2'$.
All sources are assumed to be at the same redshift. so.	All sources are assumed to be at the same redshift $z_{\rm s}$.
Whereas the number density of haloes as measured. with AZ, depends strongly on the source redshift. this dependence is quite similar in all cosmologies. except at rather low recdshifts. τον 0.6.	Whereas the number density of haloes as measured with $M_{\rm ap}$ depends strongly on the source redshift, this dependence is quite similar in all cosmologies, except at rather low redshifts, $z_{\rm s}\sim 0.6$ .
Llowever. their number density is likely to be fairly small. so that the differences seen in refsrecshilt will be very cliflicult to measure.	However, their number density is likely to be fairly small, so that the differences seen in \\ref{sredshift} will be very difficult to measure.
We therefore discard this indicator at this point.	We therefore discard this indicator at this point.
In this paper we investigated the statistics of high signal-to- noise peaks in the aperture mass map in various cosmological nmiocdels.	In this paper we investigated the statistics of high signal-to- noise peaks in the aperture mass map in various cosmological models.
We constructed the observable number of peaks in the aperture mass VCsMa8) using the Press-Schechter theory for evaluating the number clonsity of haloes anc the universal density. profile of NEW.	We constructed the observable number of peaks in the aperture mass $N(> M_{\rm ap}, \theta)$ using the Press-Schechter theory for evaluating the number density of haloes and the universal density profile of NFW.
Phe observable number density of high signal-to-noise peaks in the aperture mass map or in other words. the number density of masseselected haloes — is large in all cosmological models considered here. and range [rom ~10deg7 for a cluster normalized. Eds model to ~70deg? for a CODBE-normalized EdS model. quoted for a signal-to-noise ratio of 5.	The observable number density of high signal-to-noise peaks in the aperture mass map – or in other words, the number density of mass-selected haloes – is large in all cosmological models considered here, and range from $\sim 10 \ {\rm deg}^{-2}$ for a cluster normalized EdS model to $\sim 70 \ {\rm deg}^{-2}$ for a COBE-normalized EdS model, quoted for a signal-to-noise ratio of 5.
Even for a signal-to-noise ratio of 10. the number density of detectable haloes is about one per square degree.	Even for a signal-to-noise ratio of 10, the number density of detectable haloes is about one per square degree.
Llencee. in future wide-ield imaging surveys. such haloes will casily be found. so hat a masseselected. sample of "clusters! is within reach.	Hence, in future wide-field imaging surveys, such haloes will easily be found, so that a mass-selected sample of `clusters' is within reach.
Given that the cluster abundance has been used extensively as a cosmological probe. this mass-selectecd sample will oe extremely useful to related. observations to theoretical »ecdietions.	Given that the cluster abundance has been used extensively as a cosmological probe, this mass-selected sample will be extremely useful to related observations to theoretical predictions.
We estimated that a few square degrees of a deep wide- imaging survey will be sullicient to distinguish between some of the most popular cosmological parameter sets.	We estimated that a few square degrees of a deep wide-field imaging survey will be sufficient to distinguish between some of the most popular cosmological parameter sets.
In xuticular. cluster-normatlized low-density universes can be easily. distinguished. from a cluster-normalized EdS model. which is mainly due to the fact that in the latter. the number density of haloes at a recdshift of zi0.3. which is mainly probed by our technique. is predicted to be considerably lower than in the open mocels.	In particular, cluster-normalized low-density universes can be easily distinguished from a cluster-normalized EdS model, which is mainly due to the fact that in the latter, the number density of haloes at a redshift of $z_{\rm d}\sim 0.3$, which is mainly probed by our technique, is predicted to be considerably lower than in the open models.
Whereas our estimates on the number density of detectable haloes are. based. on. several. simplifying assumptions (e.g.. that halo number density can be obtained from the Press-Schechter theory. that the mass density is spherical and follows an NEW profile. that halos are isolated. ete.)	Whereas our estimates on the number density of detectable haloes are based on several simplifying assumptions (e.g., that halo number density can be obtained from the Press-Schechter theory, that the mass density is spherical and follows an NFW profile, that halos are isolated, etc.)
and therefore probably not very accurate. the numbers obtained should. approximately reflect. the true. situation.	and therefore probably not very accurate, the numbers obtained should approximately reflect the true situation.
In particular. the relative abundance as a function of AM, and in dependence on cosmological parameters will be the same as calculated: here.	In particular, the relative abundance as a function of $M_{\rm ap}$ and in dependence on cosmological parameters will be the same as calculated here.
For more quantitative estimates. rav-tracing calculations in a model universe obtained fron N-body simulations have to be used.	For more quantitative estimates, ray-tracing calculations in a model universe obtained from N-body simulations have to be used.
With results obtained from there. more sensitive statistics for the determination of cosmological parameters can be derived.	With results obtained from there, more sensitive statistics for the determination of cosmological parameters can be derived.
We want to thank CoS. Frenk for. providing us the routine for computing the NEW profile parameter.	We want to thank C.S. Frenk for providing us the routine for computing the NFW profile parameter.
“Phis work was supported. by the 7Sonderforschungsbereich 315-05 [fürr Astro-Teilchenphysik" der Deutschenl'orschungsgemeinschatt.	This work was supported by the “Sonderforschungsbereich 375-95 fürr Astro-Teilchenphysik" der DeutschenForschungsgemeinschaft.
TheAPEX photometry was cou»ared to the results Grom both aperture sizes.	The photometry was compared to the results from both aperture sizes.
No systematic rends with maguitucde were fouucd.	No systematic trends with magnitude were found.
Tje clilerence between tlie aud tlie 3-pixel aperture [Iuxes were cousistent with the published aperture correction for trausformatious between the 3 and 10 jxel apertures.	The difference between the and the 3-pixel aperture fluxes were consistent with the published aperture correction for transformations between the 3 and 10 pixel apertures.
Both the auc Sla\| aperture results were then compared with the large aperture values.	Both the and small aperture results were then compared with the large aperture values.
As was expected. tle sCater in the large aperture results was much larger thau he other two methods. but no systeiialles were found.	As was expected, the scatter in the large aperture results was much larger than the other two methods, but no systematics were found.
Aperture correctious were derived [or each uid. aud agreed with the cold wissicyu valies (1.12E and 1.127 at 3.6 pan and E25 pin respectively) o within154.	Aperture corrections were derived for each band, and agreed with the cold mission values (1.124 and 1.127 at 3.6 $\mu$ m and 4.5 $\mu$ m respectively) to within.
. To put the PRE photometry on the same system as the standard aperture photometry the ollowiug correction was used: where F is the flux in the staucdard system. Fpyy is the [lux output byAPEX. A is the aperture correction (1.03 and 1.02 for channels 1 aud 2 respectively) aud PPCog and PPCη are the cold and warm pixel phase corrections.	To put the PRF photometry on the same system as the standard aperture photometry the following correction was used: where $F$ is the flux in the standard system, $F_{PRF}$ is the flux output by, $A$ is the aperture correction (1.03 and 1.02 for channels 1 and 2 respectively) and $PPC_{Cold}$ and $PPC_{Warm}$ are the cold and warm pixel phase corrections.
The fluxes were converted to magnitudes using the standard zero-imagnitude flux densities of 280.9 Jv ([3.6]) and 179.7 Jv ([L5]).	The fluxes were converted to magnitudes using the standard zero–magnitude flux densities of 280.9 Jy ([3.6]) and 179.7 Jy ([4.5]).
We first refined the available periods as follows.	We first refined the available periods as follows.
Light curves were produced in the 3.6 and L2 yan bands.	Light curves were produced in the 3.6 and 4.5 $\mu$ m bands.
The data poiuts from Madoreetal.(2009).. taken as part of the SAGE project (Meixueretal.2006).. were also phliased in.	The data points from \citet{2009ApJ...695..988M}, taken as part of the SAGE project \citep{2006AJ....132.2268M}, were also phased in.
The periods given in POL were taken as initial values: these were unproved by iterating on the period aud re-phasiug the data until a Τὰ dispersion in the light curves was obtained in all available wavebands (U.B.V.I.£.4.H.WV.[3.6].[E35].[5.8].[8.0]. but note that not all wavelengths were available for all Cephlieids).	The periods given in P04 were taken as initial values; these were improved by iterating on the period and re-phasing the data until a minimum dispersion in the light curves was obtained in all available wavebands $U,B,V,R,I,J,H,K,[3.6],[4.5],[5.8],[8.0]$, but note that not all wavelengths were available for all Cepheids).
Ou average. the periods changed by less than196.	On average, the periods changed by less than.
. In some cases it was clear that the period of the Cepheid had changed iu the time interval between the optical and iufrared observations.	In some cases it was clear that the period of the Cepheid had changed in the time interval between the optical and infrared observations.
This is to be expected. as we utilized as much. archival data as possible: tie. time baseline for some Cepheids was of the order of many decades. which is comparable to the time many Cepheids have been found to siguificautly cliauge heir periods (loradiscussionofperiodchangesinsampleLMIC 2001)..	This is to be expected, as we utilized as much archival data as possible; the time baseline for some Cepheids was of the order of many decades, which is comparable to the time many Cepheids have been found to significantly change their periods \citep[for a discussion of period changes in a sample of LMC Cepheids, see][]{2001AcA....51..247P}.
The practice of fitting all wavebauds simultaneously makes it easy to find objects for which his is the case. as the period that fits the later observations will result in an incorrectly phased ight curve for the archival data. au vice versa.	The practice of fitting all wavebands simultaneously makes it easy to find objects for which this is the case, as the period that fits the later observations will result in an incorrectly phased light curve for the archival data, and vice versa.
When this was the case the IRAC and JAW data alone were used to refine the period aud to phase the data.	When this was the case the IRAC and $JHK$ data alone were used to refine the period and to phase the data.
The finally adopted periods are given iu Table 1..	The finally adopted periods are given in Table \ref{tab:mag_results}. .
We investigated several sources with the Effelsberg 100-m telescope for HCN direct f-type transitions, with no detections in G31.41+0.31, G34.26+0.15, W3(H20), and W51d.	We investigated several sources with the Effelsberg 100-m telescope for HCN direct $\ell$ -type transitions, with no detections in G31.41+0.31, G34.26+0.15, $_2$ O), and W51d.
Toward W5le, the J=12 line was detected (with half the intensity of G10.47+0.03), but the lower-J transitions are questionable (J=9 at least 10 times lower than in G10.47+0.03).	Toward W51e, the $J$ =12 line was detected (with half the intensity of G10.47+0.03), but the $J$ transitions are questionable $J$ =9 at least 10 times lower than in G10.47+0.03).
Clear detections in Galactic star-forming regions were only made in Orion-KL, G10.47+0.03, SgrB2-N, and -M. These latter three sources are shown in Fig. 8,,	Clear detections in Galactic star-forming regions were only made in Orion-KL, G10.47+0.03, SgrB2-N, and -M. These latter three sources are shown in Fig. \ref{fig:sd},
with additional IRAM 30-m data for SgrB2.	with additional IRAM 30-m data for SgrB2.
J=19 is contaminated by an Hy recombination line whose frequency corresponds to a ~20 km s! higher velocity.	$J$ =19 is contaminated by an $\gamma$ recombination line whose frequency corresponds to a $\sim$ 20 km $^{-1}$ higher velocity.
J=20 is blended with a transition of ethanol in SgrB2-N and possibly with Hey in -M. To compare our VLA data to the single-dish data, Fig.	$J$ =20 is blended with a transition of ethanol in SgrB2-N and possibly with $\gamma$ in -M. To compare our VLA data to the single-dish data, Fig.
9 shows the line integrated over the whole map of Figs. 2., 4,,	\ref{fig:lineflux} shows the line integrated over the whole map of Figs. \ref{fig:g10_maps}, \ref{fig:b2n_maps},
and 6..	and \ref{fig:b2m_maps}.
The VLA flux in G10.47+0.03 and SgrB2-N agrees well with the expectations from the Effelsberg data.	The VLA flux in G10.47+0.03 and SgrB2-N agrees well with the expectations from the Effelsberg data.
In SgrB2-M, absorption dominates at low frequencies, but in the IRAM 30-m data weak emission can be seen.	In SgrB2-M, absorption dominates at low frequencies, but in the IRAM 30-m data weak emission can be seen.
Emission also contributes to the VLA flux.	Emission also contributes to the VLA flux.
A bit puzzling is the different widths of J=9 and 10 in the Effelsberg data of SgrB2-M: It is 20 km s! in J=9 and only 6.6 km s! in J=10 (Fig. 8)).	A bit puzzling is the different widths of $J$ =9 and 10 in the Effelsberg data of SgrB2-M: It is 20 km $^{-1}$ in $J$ =9 and only 6.6 km $^{-1}$ in $J$ =10 (Fig. \ref{fig:sd}) ).
As we found no blending at the J=9 frequency, this is probably due to different ratios of the 50, 60, and 70 km s! components: For J=10, the 60 km s! component dominates.	As we found no blending at the $J$ =9 frequency, this is probably due to different ratios of the 50, 60, and 70 km $^{-1}$ components: For $J$ =10, the 60 km $^{-1}$ component dominates.
In SgrB2-N, also the higher-J transitions with level energies up to 2100 K were clearly detected.	In SgrB2-N, also the $J$ transitions with level energies up to 2100 K were clearly detected.
This allows to construct a rotation diagram (Fig. 10)),	This allows to construct a rotation diagram (Fig. \ref{fig:rd}) ),
yielding a temperature of 485+50 K and an HCN column density in a 1" source of (5.8+2)x10? cm.	yielding a temperature of $485\pm 50$ K and an HCN column density in a $''$ source of $(5.8\pm 2) \times 10^{19}$ $^{-2}$.
Note that taking absorption and optical depth into account would lead to a higher column density and a lower temperature, since both effects weaken especially the lower-J lines.	Note that taking absorption and optical depth into account would lead to a higher column density and a lower temperature, since both effects weaken especially the $J$ lines.
In an attempt to constrain the spatial structure of the hot molecular gas, we constructed radiative-transfer models that reproduce the observations.	In an attempt to constrain the spatial structure of the hot molecular gas, we constructed radiative-transfer models that reproduce the observations.
We employed a trial-and-error	We employed a trial-and-error
Fic. 1a	. .—
. The COL O spectu of PSS 2322\|1911 as measured at 22.515 GIIz at a spatial resolutiou of 3.8".	The CO 1–0 spectrum of PSS 2322+1944 as measured at 22.515 GHz at a spatial resolution of $''$.
Zero volocity is defined as the center of chamnel l. corresponding to a hehocentric redshift of 1.11956 for CO 10.	Zero velocity is defined as the center of channel 4, corresponding to a heliocentric redshift of 4.11956 for CO 1–0.
Each chaune is 83 kus ! wide. aud the Pls nolse du Czihi channel is 0.12 mJy beam.	Each channel is 83 km $^{-1}$ wide, and the rms noise in each channel is 0.12 mJy $^{-1}$.
The dashed line is a Caussian fit to the velocity xofile (see section 3).	The dashed line is a Gaussian fit to the velocity profile (see section 3).
lb.. The contour image of the average of chanucls 3. 1. and 5 from the spectitmi shown in Fieure la.	The contour image of the average of channels 3, 4, and 5 from the spectrum shown in Figure 1a.
The contour levels are: -0.21. -O.12. 0.12. 0.21. 0.36. 0.Ls. 1.60 wJv beamο,	The contour levels are: -0.24, -0.12, 0.12, 0.24, 0.36, 0.48, 0.60 mJy $^{-1}$.
The Caussian restoring CLEAN vun has FWIIM = 3.9"&3.7 "with a imnajor axis »osition angle of 167. as shown in the inset.	The Gaussian restoring CLEAN beam has FWHM = $3.9'' \times 3.7''$ with a major axis position angle of $^\circ$, as shown in the inset.
The riis voise on this image is FlyJy \|.	The rms noise on this image is $\mu$ Jy $^{-1}$.
The crosses iu his ancl subsequent niages show the position of the wo optical QSOs found by Djorgovski et al. (	The crosses in this and subsequent images show the position of the two optical QSOs found by Djorgovski et al. (
2002)1c.	2002).
. The contour image of the off-line channels (1 aud 6) from the spectrin shown in Figure 1a.	The contour image of the off-line channels (1 and 6) from the spectrum shown in Figure 1a.
The contour evels aud heau are the same as Figure 1b.	The contour levels and beam are the same as Figure 1b.
Fic. 2a.The coutour image of the average of he ou-line IF for the CO 21 observations of PSS 2322\|1911.	.—The contour image of the average of the on-line IF for the CO 2–1 observations of PSS 2322+1944.
The IF bandwidth is 50 MIIz (= 350 lau ly and is centered at 15.035 €Uz. correspoudine to a helioceutric redshitt of L1191 GIIz for CO 21.	The IF bandwidth is 50 MHz (= 350 km $^{-1}$ ), and is centered at 45.035 GHz, corresponding to a heliocentric redshift of 4.1191 GHz for CO 2–1.
The contour levels are -0.21. -0.12. (N12. 0.21. 0.36. 0.15. 1.60 wJv beamο,	The contour levels are -0.24, -0.12, 0.12, 0.24, 0.36, 0.48, 0.60 mJy $^{-1}$.
The Caussian restoring CLEAN sun has FWIIM = 1.8”SLY with a major axis xositiou angle of -28". as shown in the iuset.	The Gaussian restoring CLEAN beam has FWHM = $1.8'' \times 1.4''$ with a major axis position angle of $^\circ$, as shown in the inset.
The riis roise on this image is Ty dy \|.2b..	The rms noise on this image is $\mu$ Jy $^{-1}$.
The contour tage of the offline IF for the 11.985 11 VLA observations of 23221911.	The contour image of the off-line IF for the 44.985 GHz VLA observations of 2322+1944.
The coutour evels aud heau are the same as Figure 2a.Fi	The contour levels and beam are the same as Figure 2a. —
c3. The velocitvanteerated fiux densitics for CO enussion from PSS 232241911. as derived. from data from this paper aid from Cox et al. (	The velocity-integrated flux densities for CO emission from PSS 2322+1944, as derived from data from this paper and from Cox et al. (
2002). are show as laree filled circles with error bars.	2002), are show as large filled circles with error bars.
Also plotted are the corresponding valucs for the QSOs BRI 12020725 at 2=Lf (Opel squares). and BRI 1335ο τα =Ll (open triangles) (from Carilli et al.	Also plotted are the corresponding values for the QSOs BRI 1202–0725 at $z = 4.7$ (open squares), and BRI 1335–0417 at $z = 4.4$ (open triangles) (from Carilli et al.
2002).	2002).
The solid truges are the data for the CO ladder for the iuteerate( emission from the Alilkv Wav ciss inside the solar radius (excluding the Galactic center) as seen by CODE (Fixsen. Deunett. Mather 1999).	The solid triangles are the data for the CO ladder for the integrated emission from the Milky Way disk inside the solar radius (excluding the Galactic center) as seen by COBE (Fixsen, Bennett, Mather 1999).
The open cireles are the results for the μίαν]1rst nucleus of M82 (Caissten oet al.	The open circles are the results for the starburst nucleus of M82 (Güssten et al.
1993: Mao e al.	1993; Mao et al.
2000).	2000).
The oxinate is the velocity-iutegerated line flux density.	The ordinate is the velocity-integrated line flux density.
The CO(5-D) line has beeu detected iu all he sources, and t10 velocity-integrated line flux densities for the other trausifious in cach source lave all been normalized by the corresponding CO(S-1) values. except for the Milky Way. for which the values ave normalized at CO(3-2).	The CO(5-4) line has been detected in all the sources, and the velocity-integrated line flux densities for the other transitions in each source have all been normalized by the corresponding CO(5-4) values, except for the Milky Way, for which the values are normalized at CO(3-2).
The short dash line shows an LVC model with Tig, = {τν aud u(IT») Ξ ον10 5m	The short dash line shows an LVG model with $_{\rm kin}$ = 47 K and $_2$ ) = $5 \times 10^3$ $^{-3}$ .

]t is important to have in münd that the process of structure. formation of the Universe. ds. a very complicated one.

It is important to have in mind that the process of structure formation of the Universe is a very complicated one.

For the canonical model of cosmology. namely primordial Gaussian random perturbation field. the formation of voids and clusters of galaxies is part of one process of clustering.

For the canonical model of cosmology, namely primordial Gaussian random perturbation field, the formation of voids and clusters of galaxies is part of one process of clustering.

The formation of regions with positive density contrasts. like clusters of galaxies. contributes to the erowth of underdense regions and. vice versa.

The formation of regions with positive density contrasts, like clusters of galaxies, contributes to the growth of underdense regions and vice versa.

The present study only. considers a small part. of the very complicated: structure formation. process.

The present study only considers a small part of the very complicated structure formation process.

We study the evolution of negative density perturbations taking into account all relevant physical processes acting on them.

Llowever. the assessment of the role of these very physical processes not considered in the literature. besides including the presence of non-barvonic dark matter ancl cosmological constant. are the maincontributions of the present studs.

However, the assessment of the role of these very physical processes not considered in the literature, besides including the presence of non-baryonic dark matter and cosmological constant, are the maincontributions of the present study.

idoes per square degree with aperture mass greater than Aly=0.04 and AL,=0.08 for the filter scale @=2’ or the five cosmological models. using two different redshift intervals. zO.15.04]) and. z0.4.1].

haloes per square degree with aperture mass greater than $M_{\rm ap}=0.04$ and $M_{\rm ap}=0.08$ for the filter scale $\theta=2'$ for the five cosmological models, using two different redshift intervals, $z\in[0.15,0.4]$ and $z\in[0.4,1]$.

By. comparing he halo densities in the different recshift intervals for the various cosmologies in refndachx. and 3... we expect the Iargest dillerences between he cosmological models for AZ,=0.08.

By comparing the halo densities in the different redshift intervals for the various cosmologies in \\ref{ndachx} and \ref{ndach1}, we expect the largest differences between the cosmological models for $M_{\rm ap}=0.08$.

The reason for this is the stronger evolution for the rich cluster mass function which corresponds to large values of the aperture mass (sce Figure 1)).

The reason for this is the stronger evolution for the rich cluster mass function which corresponds to large values of the aperture mass (see Figure \ref{map}) ).

Whereas the EdS(0.6.0.25) and EcdS(1.0.25) models are again very ilferent. from the other three. the use of redshift information greatly helps to distinguish the EdS(0.6.0.5) model from the two low-clensity niocels.

Whereas the EdS(0.6,0.25) and EdS(1,0.25) models are again very different from the other three, the use of redshift information greatly helps to distinguish the EdS(0.6,0.5) model from the two low-density models.

For the latter. a survey area of less than 3 deg? would be sullicient.

For the latter, a survey area of less than 3 ${\rm deg}^2$ would be sufficient.

One might think of another way to obtain reclshilt information. namely to use source galaxies at. cillerent redshilts (clistinguishecd. say. by photometric redshift estimates).

One might think of another way to obtain redshift information, namely to use source galaxies at different redshifts (distinguished, say, by photometric redshift estimates).

To investigate this effect. we have plottec in Figure 7. the dependence of the number of haloes on the redshift of the sources for AL,=0.04 and 8=2’.

To investigate this effect, we have plotted in Figure \ref{sredshift} the dependence of the number of haloes on the redshift of the sources for $M_{\rm ap}=0.04$ and $\theta=2'$.

All sources are assumed to be at the same redshift. so.

All sources are assumed to be at the same redshift $z_{\rm s}$.

Whereas the number density of haloes as measured. with AZ, depends strongly on the source redshift. this dependence is quite similar in all cosmologies. except at rather low recdshifts. τον 0.6.

Whereas the number density of haloes as measured with $M_{\rm ap}$ depends strongly on the source redshift, this dependence is quite similar in all cosmologies, except at rather low redshifts, $z_{\rm s}\sim 0.6$ .

Llowever. their number density is likely to be fairly small. so that the differences seen in refsrecshilt will be very cliflicult to measure.

However, their number density is likely to be fairly small, so that the differences seen in \\ref{sredshift} will be very difficult to measure.

We therefore discard this indicator at this point.

In this paper we investigated the statistics of high signal-to- noise peaks in the aperture mass map in various cosmological nmiocdels.

In this paper we investigated the statistics of high signal-to- noise peaks in the aperture mass map in various cosmological models.

We constructed the observable number of peaks in the aperture mass VCsMa8) using the Press-Schechter theory for evaluating the number clonsity of haloes anc the universal density. profile of NEW.

We constructed the observable number of peaks in the aperture mass $N(> M_{\rm ap}, \theta)$ using the Press-Schechter theory for evaluating the number density of haloes and the universal density profile of NFW.

Phe observable number density of high signal-to-noise peaks in the aperture mass map or in other words. the number density of masseselected haloes — is large in all cosmological models considered here. and range [rom ~10deg7 for a cluster normalized. Eds model to ~70deg? for a CODBE-normalized EdS model. quoted for a signal-to-noise ratio of 5.

The observable number density of high signal-to-noise peaks in the aperture mass map – or in other words, the number density of mass-selected haloes – is large in all cosmological models considered here, and range from $\sim 10 \ {\rm deg}^{-2}$ for a cluster normalized EdS model to $\sim 70 \ {\rm deg}^{-2}$ for a COBE-normalized EdS model, quoted for a signal-to-noise ratio of 5.

Even for a signal-to-noise ratio of 10. the number density of detectable haloes is about one per square degree.

Even for a signal-to-noise ratio of 10, the number density of detectable haloes is about one per square degree.

Llencee. in future wide-ield imaging surveys. such haloes will casily be found. so hat a masseselected. sample of "clusters! is within reach.

Hence, in future wide-field imaging surveys, such haloes will easily be found, so that a mass-selected sample of `clusters' is within reach.

Given that the cluster abundance has been used extensively as a cosmological probe. this mass-selectecd sample will oe extremely useful to related. observations to theoretical »ecdietions.

Given that the cluster abundance has been used extensively as a cosmological probe, this mass-selected sample will be extremely useful to related observations to theoretical predictions.

We estimated that a few square degrees of a deep wide- imaging survey will be sullicient to distinguish between some of the most popular cosmological parameter sets.

We estimated that a few square degrees of a deep wide-field imaging survey will be sufficient to distinguish between some of the most popular cosmological parameter sets.

In xuticular. cluster-normatlized low-density universes can be easily. distinguished. from a cluster-normalized EdS model. which is mainly due to the fact that in the latter. the number density of haloes at a recdshift of zi0.3. which is mainly probed by our technique. is predicted to be considerably lower than in the open mocels.

In particular, cluster-normalized low-density universes can be easily distinguished from a cluster-normalized EdS model, which is mainly due to the fact that in the latter, the number density of haloes at a redshift of $z_{\rm d}\sim 0.3$, which is mainly probed by our technique, is predicted to be considerably lower than in the open models.

Whereas our estimates on the number density of detectable haloes are. based. on. several. simplifying assumptions (e.g.. that halo number density can be obtained from the Press-Schechter theory. that the mass density is spherical and follows an NEW profile. that halos are isolated. ete.)

Whereas our estimates on the number density of detectable haloes are based on several simplifying assumptions (e.g., that halo number density can be obtained from the Press-Schechter theory, that the mass density is spherical and follows an NFW profile, that halos are isolated, etc.)

and therefore probably not very accurate. the numbers obtained should. approximately reflect. the true. situation.

and therefore probably not very accurate, the numbers obtained should approximately reflect the true situation.

In particular. the relative abundance as a function of AM, and in dependence on cosmological parameters will be the same as calculated: here.

In particular, the relative abundance as a function of $M_{\rm ap}$ and in dependence on cosmological parameters will be the same as calculated here.

For more quantitative estimates. rav-tracing calculations in a model universe obtained fron N-body simulations have to be used.

For more quantitative estimates, ray-tracing calculations in a model universe obtained from N-body simulations have to be used.

With results obtained from there. more sensitive statistics for the determination of cosmological parameters can be derived.

With results obtained from there, more sensitive statistics for the determination of cosmological parameters can be derived.

We want to thank CoS. Frenk for. providing us the routine for computing the NEW profile parameter.

We want to thank C.S. Frenk for providing us the routine for computing the NFW profile parameter.

“Phis work was supported. by the 7Sonderforschungsbereich 315-05 [fürr Astro-Teilchenphysik" der Deutschenl'orschungsgemeinschatt.

This work was supported by the “Sonderforschungsbereich 375-95 fürr Astro-Teilchenphysik" der DeutschenForschungsgemeinschaft.

TheAPEX photometry was cou»ared to the results Grom both aperture sizes.

The photometry was compared to the results from both aperture sizes.

No systematic rends with maguitucde were fouucd.

No systematic trends with magnitude were found.

Tje clilerence between tlie aud tlie 3-pixel aperture [Iuxes were cousistent with the published aperture correction for trausformatious between the 3 and 10 jxel apertures.

The difference between the and the 3-pixel aperture fluxes were consistent with the published aperture correction for transformations between the 3 and 10 pixel apertures.

Both the auc Sla| aperture results were then compared with the large aperture values.

Both the and small aperture results were then compared with the large aperture values.

As was expected. tle sCater in the large aperture results was much larger thau he other two methods. but no systeiialles were found.

As was expected, the scatter in the large aperture results was much larger than the other two methods, but no systematics were found.

Aperture correctious were derived [or each uid. aud agreed with the cold wissicyu valies (1.12E and 1.127 at 3.6 pan and E25 pin respectively) o within154.

Aperture corrections were derived for each band, and agreed with the cold mission values (1.124 and 1.127 at 3.6 $\mu$ m and 4.5 $\mu$ m respectively) to within.

. To put the PRE photometry on the same system as the standard aperture photometry the ollowiug correction was used: where F is the flux in the staucdard system. Fpyy is the [lux output byAPEX. A is the aperture correction (1.03 and 1.02 for channels 1 aud 2 respectively) aud PPCog and PPCη are the cold and warm pixel phase corrections.

To put the PRF photometry on the same system as the standard aperture photometry the following correction was used: where $F$ is the flux in the standard system, $F_{PRF}$ is the flux output by, $A$ is the aperture correction (1.03 and 1.02 for channels 1 and 2 respectively) and $PPC_{Cold}$ and $PPC_{Warm}$ are the cold and warm pixel phase corrections.

The fluxes were converted to magnitudes using the standard zero-imagnitude flux densities of 280.9 Jv ([3.6]) and 179.7 Jv ([L5]).

The fluxes were converted to magnitudes using the standard zero–magnitude flux densities of 280.9 Jy ([3.6]) and 179.7 Jy ([4.5]).

We first refined the available periods as follows.

Light curves were produced in the 3.6 and L2 yan bands.

Light curves were produced in the 3.6 and 4.5 $\mu$ m bands.

The data poiuts from Madoreetal.(2009).. taken as part of the SAGE project (Meixueretal.2006).. were also phliased in.

The data points from \citet{2009ApJ...695..988M}, taken as part of the SAGE project \citep{2006AJ....132.2268M}, were also phased in.

The periods given in POL were taken as initial values: these were unproved by iterating on the period aud re-phasiug the data until a Τὰ dispersion in the light curves was obtained in all available wavebands (U.B.V.I.£.4.H.WV.[3.6].[E35].[5.8].[8.0]. but note that not all wavelengths were available for all Cephlieids).

The periods given in P04 were taken as initial values; these were improved by iterating on the period and re-phasing the data until a minimum dispersion in the light curves was obtained in all available wavebands $U,B,V,R,I,J,H,K,[3.6],[4.5],[5.8],[8.0]$, but note that not all wavelengths were available for all Cepheids).

Ou average. the periods changed by less than196.

On average, the periods changed by less than.

. In some cases it was clear that the period of the Cepheid had changed iu the time interval between the optical and iufrared observations.

In some cases it was clear that the period of the Cepheid had changed in the time interval between the optical and infrared observations.

This is to be expected. as we utilized as much. archival data as possible: tie. time baseline for some Cepheids was of the order of many decades. which is comparable to the time many Cepheids have been found to siguificautly cliauge heir periods (loradiscussionofperiodchangesinsampleLMIC 2001)..

This is to be expected, as we utilized as much archival data as possible; the time baseline for some Cepheids was of the order of many decades, which is comparable to the time many Cepheids have been found to significantly change their periods \citep[for a discussion of period changes in a sample of LMC Cepheids, see][]{2001AcA....51..247P}.

The practice of fitting all wavebauds simultaneously makes it easy to find objects for which his is the case. as the period that fits the later observations will result in an incorrectly phased ight curve for the archival data. au vice versa.

The practice of fitting all wavebands simultaneously makes it easy to find objects for which this is the case, as the period that fits the later observations will result in an incorrectly phased light curve for the archival data, and vice versa.

When this was the case the IRAC and JAW data alone were used to refine the period aud to phase the data.

When this was the case the IRAC and $JHK$ data alone were used to refine the period and to phase the data.

The finally adopted periods are given iu Table 1..

The finally adopted periods are given in Table \ref{tab:mag_results}. .

We investigated several sources with the Effelsberg 100-m telescope for HCN direct f-type transitions, with no detections in G31.41+0.31, G34.26+0.15, W3(H20), and W51d.

We investigated several sources with the Effelsberg 100-m telescope for HCN direct $\ell$ -type transitions, with no detections in G31.41+0.31, G34.26+0.15, $_2$ O), and W51d.

Toward W5le, the J=12 line was detected (with half the intensity of G10.47+0.03), but the lower-J transitions are questionable (J=9 at least 10 times lower than in G10.47+0.03).

Toward W51e, the $J$ =12 line was detected (with half the intensity of G10.47+0.03), but the $J$ transitions are questionable $J$ =9 at least 10 times lower than in G10.47+0.03).

Clear detections in Galactic star-forming regions were only made in Orion-KL, G10.47+0.03, SgrB2-N, and -M. These latter three sources are shown in Fig. 8,,

Clear detections in Galactic star-forming regions were only made in Orion-KL, G10.47+0.03, SgrB2-N, and -M. These latter three sources are shown in Fig. \ref{fig:sd},

with additional IRAM 30-m data for SgrB2.

J=19 is contaminated by an Hy recombination line whose frequency corresponds to a ~20 km s! higher velocity.

$J$ =19 is contaminated by an $\gamma$ recombination line whose frequency corresponds to a $\sim$ 20 km $^{-1}$ higher velocity.

J=20 is blended with a transition of ethanol in SgrB2-N and possibly with Hey in -M. To compare our VLA data to the single-dish data, Fig.

$J$ =20 is blended with a transition of ethanol in SgrB2-N and possibly with $\gamma$ in -M. To compare our VLA data to the single-dish data, Fig.

9 shows the line integrated over the whole map of Figs. 2., 4,,

\ref{fig:lineflux} shows the line integrated over the whole map of Figs. \ref{fig:g10_maps}, \ref{fig:b2n_maps},

and 6..

and \ref{fig:b2m_maps}.

The VLA flux in G10.47+0.03 and SgrB2-N agrees well with the expectations from the Effelsberg data.

In SgrB2-M, absorption dominates at low frequencies, but in the IRAM 30-m data weak emission can be seen.

Emission also contributes to the VLA flux.

A bit puzzling is the different widths of J=9 and 10 in the Effelsberg data of SgrB2-M: It is 20 km s! in J=9 and only 6.6 km s! in J=10 (Fig. 8)).

A bit puzzling is the different widths of $J$ =9 and 10 in the Effelsberg data of SgrB2-M: It is 20 km $^{-1}$ in $J$ =9 and only 6.6 km $^{-1}$ in $J$ =10 (Fig. \ref{fig:sd}) ).

As we found no blending at the J=9 frequency, this is probably due to different ratios of the 50, 60, and 70 km s! components: For J=10, the 60 km s! component dominates.

As we found no blending at the $J$ =9 frequency, this is probably due to different ratios of the 50, 60, and 70 km $^{-1}$ components: For $J$ =10, the 60 km $^{-1}$ component dominates.

In SgrB2-N, also the higher-J transitions with level energies up to 2100 K were clearly detected.

In SgrB2-N, also the $J$ transitions with level energies up to 2100 K were clearly detected.

This allows to construct a rotation diagram (Fig. 10)),

This allows to construct a rotation diagram (Fig. \ref{fig:rd}) ),

yielding a temperature of 485+50 K and an HCN column density in a 1" source of (5.8+2)x10? cm.

yielding a temperature of $485\pm 50$ K and an HCN column density in a $''$ source of $(5.8\pm 2) \times 10^{19}$ $^{-2}$.

Note that taking absorption and optical depth into account would lead to a higher column density and a lower temperature, since both effects weaken especially the lower-J lines.

Note that taking absorption and optical depth into account would lead to a higher column density and a lower temperature, since both effects weaken especially the $J$ lines.

In an attempt to constrain the spatial structure of the hot molecular gas, we constructed radiative-transfer models that reproduce the observations.

We employed a trial-and-error

Fic. 1a

. .—

. The COL O spectu of PSS 2322|1911 as measured at 22.515 GIIz at a spatial resolutiou of 3.8".

The CO 1–0 spectrum of PSS 2322+1944 as measured at 22.515 GHz at a spatial resolution of $''$.

Zero volocity is defined as the center of chamnel l. corresponding to a hehocentric redshift of 1.11956 for CO 10.

Zero velocity is defined as the center of channel 4, corresponding to a heliocentric redshift of 4.11956 for CO 1–0.

Each chaune is 83 kus ! wide. aud the Pls nolse du Czihi channel is 0.12 mJy beam.

Each channel is 83 km $^{-1}$ wide, and the rms noise in each channel is 0.12 mJy $^{-1}$.

The dashed line is a Caussian fit to the velocity xofile (see section 3).

The dashed line is a Gaussian fit to the velocity profile (see section 3).

lb.. The contour image of the average of chanucls 3. 1. and 5 from the spectitmi shown in Fieure la.

The contour image of the average of channels 3, 4, and 5 from the spectrum shown in Figure 1a.

The contour levels are: -0.21. -O.12. 0.12. 0.21. 0.36. 0.Ls. 1.60 wJv beamο,

The contour levels are: -0.24, -0.12, 0.12, 0.24, 0.36, 0.48, 0.60 mJy $^{-1}$.

The Caussian restoring CLEAN vun has FWIIM = 3.9"&3.7 "with a imnajor axis »osition angle of 167. as shown in the inset.

The Gaussian restoring CLEAN beam has FWHM = $3.9'' \times 3.7''$ with a major axis position angle of $^\circ$, as shown in the inset.

The riis voise on this image is FlyJy |.

The rms noise on this image is $\mu$ Jy $^{-1}$.

The crosses iu his ancl subsequent niages show the position of the wo optical QSOs found by Djorgovski et al. (

The crosses in this and subsequent images show the position of the two optical QSOs found by Djorgovski et al. (

2002)1c.

2002).

. The contour image of the off-line channels (1 aud 6) from the spectrin shown in Figure 1a.

The contour image of the off-line channels (1 and 6) from the spectrum shown in Figure 1a.

The contour evels aud heau are the same as Figure 1b.

The contour levels and beam are the same as Figure 1b.

Fic. 2a.The coutour image of the average of he ou-line IF for the CO 21 observations of PSS 2322|1911.

.—The contour image of the average of the on-line IF for the CO 2–1 observations of PSS 2322+1944.

The IF bandwidth is 50 MIIz (= 350 lau ly and is centered at 15.035 €Uz. correspoudine to a helioceutric redshitt of L1191 GIIz for CO 21.

The IF bandwidth is 50 MHz (= 350 km $^{-1}$ ), and is centered at 45.035 GHz, corresponding to a heliocentric redshift of 4.1191 GHz for CO 2–1.

The contour levels are -0.21. -0.12. (N12. 0.21. 0.36. 0.15. 1.60 wJv beamο,

The contour levels are -0.24, -0.12, 0.12, 0.24, 0.36, 0.48, 0.60 mJy $^{-1}$.

The Caussian restoring CLEAN sun has FWIIM = 1.8”SLY with a major axis xositiou angle of -28". as shown in the iuset.

The Gaussian restoring CLEAN beam has FWHM = $1.8'' \times 1.4''$ with a major axis position angle of $^\circ$, as shown in the inset.

The riis roise on this image is Ty dy |.2b..

The rms noise on this image is $\mu$ Jy $^{-1}$.

The contour tage of the offline IF for the 11.985 11 VLA observations of 23221911.

The contour image of the off-line IF for the 44.985 GHz VLA observations of 2322+1944.

The coutour evels aud heau are the same as Figure 2a.Fi

The contour levels and beam are the same as Figure 2a. —

c3. The velocitvanteerated fiux densitics for CO enussion from PSS 232241911. as derived. from data from this paper aid from Cox et al. (

The velocity-integrated flux densities for CO emission from PSS 2322+1944, as derived from data from this paper and from Cox et al. (

2002). are show as laree filled circles with error bars.

2002), are show as large filled circles with error bars.

Also plotted are the corresponding valucs for the QSOs BRI 12020725 at 2=Lf (Opel squares). and BRI 1335ο τα =Ll (open triangles) (from Carilli et al.

Also plotted are the corresponding values for the QSOs BRI 1202–0725 at $z = 4.7$ (open squares), and BRI 1335–0417 at $z = 4.4$ (open triangles) (from Carilli et al.

2002).

The solid truges are the data for the CO ladder for the iuteerate( emission from the Alilkv Wav ciss inside the solar radius (excluding the Galactic center) as seen by CODE (Fixsen. Deunett. Mather 1999).

The solid triangles are the data for the CO ladder for the integrated emission from the Milky Way disk inside the solar radius (excluding the Galactic center) as seen by COBE (Fixsen, Bennett, Mather 1999).

The open cireles are the results for the μίαν]1rst nucleus of M82 (Caissten oet al.

The open circles are the results for the starburst nucleus of M82 (Güssten et al.

1993: Mao e al.

1993; Mao et al.

2000).

The oxinate is the velocity-iutegerated line flux density.

The ordinate is the velocity-integrated line flux density.

The CO(5-D) line has beeu detected iu all he sources, and t10 velocity-integrated line flux densities for the other trausifious in cach source lave all been normalized by the corresponding CO(S-1) values. except for the Milky Way. for which the values ave normalized at CO(3-2).

The CO(5-4) line has been detected in all the sources, and the velocity-integrated line flux densities for the other transitions in each source have all been normalized by the corresponding CO(5-4) values, except for the Milky Way, for which the values are normalized at CO(3-2).

The short dash line shows an LVC model with Tig, = {τν aud u(IT») Ξ ον10 5m

The short dash line shows an LVG model with $_{\rm kin}$ = 47 K and $_2$ ) = $5 \times 10^3$ $^{-3}$ .