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RL sources show Baler lines with large red aud blue profile shifts. | RL sources show Balmer lines with large red and blue profile shifts. |
RL sources show no evidence for a soft N-rav excess which may refiect a real abseuce of this component or that it is dominated by a relatively hard component related to Nav. cussion related to the radio loudness. | RL sources show no evidence for a soft X-ray excess which may reflect a real absence of this component or that it is dominated by a relatively hard component related to X-ray emission related to the radio loudness. |
We favor the former interpretation because most RO population D sources also show a soft Χταν deficit. | We favor the former interpretation because most RQ population B sources also show a soft X-ray deficit. |
The absence of a soft N-rav component iu RL|RQ pop D sources coupled with the absence of the CIV blueshift are cousisteut with weak or abseut disk structure. | The absence of a soft X-ray component in RL+RQ pop B sources coupled with the absence of the CIV blueshift are consistent with weak or absent disk structure. |
The broad lines may arise in a biconical structure in many or all of these sources (Marzgiani ct al. | The broad lines may arise in a biconical structure in many or all of these sources (Marziani et al. |
1993: Suleutic et al. | 1993; Sulentic et al. |
1995). | 1995). |
Exideuce for à bicoue oriein of the BLR includes: a) sources with double peaked broad lines (c.g. Sulentic et al. | Evidence for a bicone origin of the BLR includes: a) sources with double peaked broad lines (e.g. Sulentic et al. |
1995: Exacleous Halpern 1998). b) sources with single peaked broad lines showing both large red aud blue dcisplaceimoeuts (οιο, Marzio et al. | 1995; Eracleous Halpern 1998), b) sources with single peaked broad lines showing both large red and blue displacements (e.g. Marziani et al. |
1983: Παρα et al. | 1983; Halpern et al. |
1998). ο) sources with a transicut DER. (c.g. Storchi-Beremanu ct al. | 1998), c) sources with a transient BLR (e.g. Storchi-Bergmann et al. |
1993: Ulrich 2000) aud spectropolariuuctry of a) aud b) sources (Corbett et al. | 1993; Ulrich 2000) and spectropolarimetry of a) and b) sources (Corbett et al. |
L998). | 1998). |
The El parameter space is measuring aspects of both the ecometry aud kinematic of the broad line region in ACN. | The E1 parameter space is measuring aspects of both the geometry and kinematic of the broad line region in AGN. |
The streneth of the correlations and reasonable orthogonality of the parameters sugecst hat a better diagnostic space for ACN is uulikelv to be found. | The strength of the correlations and reasonable orthogonality of the parameters suggest that a better diagnostic space for AGN is unlikely to be found. |
Even in the preliminary preseutation. El allows us to resolve some ACN conundrums. | Even in the preliminary presentation, E1 allows us to resolve some AGN conundrums. |
1) RL sources ave found to be an ACN population with undamental eeometrical and imematic difference from the RQ majority. | 1) RL sources are found to be an AGN population with fundamental geometrical and kinematic difference from the RQ majority. |
2) NLSvl are found to be an extremmun of the RQ population rather than a pathological or disjoint ACN population. | 2) NLSy1 are found to be an extremum of the RQ population rather than a pathological or disjoint AGN population. |
3) If the width of he Bahluner lines bear auy signature of source orientation then some NLSvil (at least. I Zw 1) are secu at or rear pole-on (face-on accretion disk) orientation. while some BAL QSOs are likely to be misaligned NLSv1. | 3) If the width of the Balmer lines bear any signature of source orientation then some NLSy1 (at least, I Zw 1) are seen at or near pole-on (face-on accretion disk) orientation, while some BAL QSOs are likely to be misaligned NLSy1. |
1) Population D. RO quasars with Balmer lines broader than z £000 παν represent a distinct class (RL pre/post cursors) from narrower population A RQ sources. | 4) Population B RQ quasars with Balmer lines broader than $\approx$ 4000 may represent a distinct class (RL pre/post cursors) from narrower population A RQ sources. |
5) line shifts indicate the preseuce in Pop. | 5) line shifts indicate the presence in Pop. |
A (and possibly Pop. | A (and possibly Pop. |
B) RQ sources of an ubiquitous disk outflow/wind. | B) RQ sources of an ubiquitous disk outflow/wind. |
PAL acknowledges financial support from the Italian Ministry of University aud Scicutific and Technological Research (MURST) through erant Cofiu 98-02-32 and from the DA-UNAM where part of this | PM acknowledges financial support from the Italian Ministry of University and Scientific and Technological Research (MURST) through grant Cofin 98-02-32 and from the IA-UNAM where part of this |
Tn Coorav (199958: hereafter C99). we studied the former case aud suggested the possibility of coustraiming cosmoloeical parameters based on observed statistics of leusecl sources at 850 gan towards a sample of galaxy clusters. | In Cooray (1999a; hereafter C99), we studied the former case and suggested the possibility of constraining cosmological parameters based on observed statistics of lensed sources at 850 $\mu$ m towards a sample of galaxy clusters. |
This approach is quite simular to the one taken in literature to coustrain the prescut day value of the cosmological constaut based onu leused source statistics. stich as quasars (οιοι, Iochauek 1996: Chiba Yoshi 1997). radio sources (e.g. Falco et al. | This approach is quite similar to the one taken in literature to constrain the present day value of the cosmological constant based on lensed source statistics, such as quasars (e.g., Kochanek 1996; Chiba Yoshi 1997), radio sources (e.g., Falco et al. |
1998: Coorav 19995) and Iuniuous optical arcs towards clusters (e.c... Bartehuann et al. | 1998; Cooray 1999b) and luminous optical arcs towards clusters (e.g., Bartelmann et al. |
1998: €99). | 1998; C99). |
Blain (1997) studied eravitational lensing at παπα waveleusths. iuclhudiug statistics of leused subuuu sources due to foreeround ealaxies. | Blain (1997) studied gravitational lensing at submm wavelengths, including statistics of lensed submm sources due to foreground galaxies. |
Tere. we consider the possibility of coustraimine the backeround source redshift distribution based on kuown properties of a sample of subinm sources gravitationally lensed by cluster) potentials (S98) aud an assumed cosiuological model. | Here, we consider the possibility of constraining the background source redshift distribution based on known properties of a sample of submm sources gravitationally lensed by cluster potentials (S98) and an assumed cosmological model. |
To obtain information on the uuleused sources down to the same flux level. we use the source counts frou, Darger et al. ( | To obtain information on the unlensed sources down to the same flux level, we use the source counts from Barger et al. ( |
1998). Eales et al. ( | 1998), Eales et al. ( |
1998). IIollaud et al. ( | 1998), Holland et al. ( |
1998). IIughes et al. ( | 1998), Hughes et al. ( |
1998) aud Sunuül et al. ( | 1998) and Smail et al. ( |
1997. 19958b). | 1997, 1998b). |
In 2 we discuss our calculation ancl its iuputs. | In 2 we discuss our calculation and its inputs. |
in 3 we present our resulting constraints on the redshift distribution of πιά sources aud discuss our constraiuts in the context of current studies on the stb sources and their coutribution to the starformation listory of the universe. | in 3 we present our resulting constraints on the redshift distribution of submm sources and discuss our constraints in the context of current studies on the submm sources and their contribution to the starformation history of the universe. |
We follow the conventions that the ITubble coustant. ZIy. is LOO Fin + 1. the present mean density in the universe in uuits of the closure deusitv is Q,,. and the preseut normalized cosmological coustaut is Q4. | We follow the conventions that the Hubble constant, $H_0$ , is $h$ km $^{-1}$ $^{-1}$, the present mean density in the universe in units of the closure density is $\Omega_m$, and the present normalized cosmological constant is $\Omega_\Lambda$. |
Iu a flat universe. O,,|O4=I. | In a flat universe, $\Omega_m+\Omega_\Lambda=1$. |
Our calculation follows C99 in which we calculated the expected nuuber of huninous optical arcs. radio sources id subuun galaxies towards ealaxy clusters as a function : cosinologw. | Our calculation follows C99 in which we calculated the expected number of luminous optical arcs, radio sources and submm galaxies towards galaxy clusters as a function of cosmology. |
Here. we prescribe the foreground cluster »pulatiou to be similar to what was observed by SOS xd niodel them as singular isothermal spheres (SIS) with velocitydispersion e. | Here, we prescribe the foreground cluster population to be similar to what was observed by S98 and model them as singular isothermal spheres (SIS) with velocity dispersion $\sigma$. |
Iu general. SIS models underestimate ie uunuber of lensed sources. when compared to complex cluster potentials with substructure (0.8... Déózzecour 1998). | In general, SIS models underestimate the number of lensed sources, when compared to complex cluster potentials with substructure (e.g., Bézzecourt 1998). |
This leads to a systematically lower nuuber of cused sources than expected frou true complex potentials iid a higher upper luit ou the redshift distribution of vackeround sources. | This leads to a systematically lower number of lensed sources than expected from true complex potentials and a higher upper limit on the redshift distribution of background sources. |
lu order to evaluate distances. we use the analvtica filled-beam approximation (see. e.g. Fukugita et al. | In order to evaluate distances, we use the analytical filled-beam approximation (see, e.g., Fukugita et al. |
1992) and calculate the probabilitv. p(z.9,,.04). for a source at redshift of + to be strongly lensed given a set of cosmoloeical parameters O,, aud O4. | 1992) and calculate the probability, $p(z,\Omega_m,\Omega_\Lambda)$, for a source at redshift of $z$ to be strongly lensed given a set of cosmological parameters $\Omega_m$ and $\Omega_\Lambda$. |
Following C99 (sce. also. Cooray. QuashnockMiller 1999 and Holz. Miller Quashnock 1999) the umuber of expected leused sources. AN. towards the survey volume contaiuiug foreerouud lensing clusters is: where C(i is the redshitt distribution of sili sources such that C(:) is the fraction of sources with redshifts less than 2. Bf.:) is the magnification bias for stb sources at redshift + with observed flax density at 850 4204 of f (sec. IKoclhanek 1991). aud. £(:) is the effectiveness of clusters at redshifts :; in producing lensed sources. | Following C99 (see, also, Cooray, QuashnockMiller 1999 and Holz, Miller Quashnock 1999) the number of expected lensed sources, $\bar N$, towards the survey volume containing foreground lensing clusters is: where $C(z)$ is the redshift distribution of submm sources such that $C(z)$ is the fraction of sources with redshifts less than $z$, $B(f,z)$ is the magnification bias for submm sources at redshift $z$ with observed flux density at 850 $\mu$ m of $f$ (see, Kochanek 1991), and $F(z_l)$ is the effectiveness of clusters at redshifts $z_l$ in producing lensed sources. |
This nondimensional parameter can be written as (Turner. Cott 1981): Tere. Ry=effly. aud οτε) is the uuuber deusitv of clusters with the velocity dispersion σ at redshift τι | This nondimensional parameter can be written as (Turner, Gott 1984): Here, $R_0=c/H_0$, and $n(z_l)$ is the number density of clusters with the velocity dispersion $\sigma$ at redshift $z_l$. |
Iu general. detailed knowledge either on the Iuninositv function or the flux distribution is required to calculate the magnification bias. | In general, detailed knowledge either on the luminosity function or the flux distribution is required to calculate the magnification bias. |
However. both these quautities are currently not known for the subnuu source sample. | However, both these quantities are currently not known for the submm source sample. |
Tustead of individual maguificatiou biases. we use cient estimates on the ΠΠ sources wmmbercounts to obtain an average value. | Instead of individual magnification biases, we use current estimates on the submm sources numbercounts to obtain an average value. |
Ifthe number counts of uuleused sources. Hy. With flux densities ereater than 5, towards a eiven area can be written as n4;X557. then magnification duc to gravitational lensing by an amplification <A inodifies the counts as: where 0; ds the lensed source counts. | If the number counts of unlensed sources, $n_{ul}$, with flux densities greater than $S_\nu$ towards a given area can be written as $n_{ul} \propto S_\nu^\alpha$, then magnification due to gravitational lensing by an amplification $A$ modifies the counts as: where $n_l$ is the lensed source counts. |
The average magnification bias is simply the ratio of leused to uuleused counts down to a flux density S,: Under the SIS scenario. the probability distribution for auplifications is PCA)=2/64Ly. aud the αι auuplifications is Aya,=2. | The average magnification bias is simply the ratio of lensed to unlensed counts down to a flux density $S_\nu$: Under the SIS scenario, the probability distribution for amplifications is $P(A) = 2/(A-1)^3$, and the minimum amplifications is $A_{\rm min}=2$. |
The average amplification for a xuuple of lensed sources is 3. | The average amplification for a sample of lensed sources is 3. |
This average value is consistent with the distribution of amplifications for the subi sources based ou detailed modeling of iudividual cluster potentials: 1.359 (S98). | This average value is consistent with the distribution of amplifications for the submm sources based on detailed modeling of individual cluster potentials: $1.3^{+5.0}_{-0.5}$ (S98). |
Since none of the observed leused sources are heavily amplified due to foreerouud potentials and that the amplification distribution is compatible with the SIS average. our use of SIS model to describe foreground clusters should not affect the results ercatly, | Since none of the observed lensed sources are heavily amplified due to foreground potentials and that the amplification distribution is compatible with the SIS average, our use of SIS model to describe foreground clusters should not affect the results greatly. |
Following Simail et al. ( | Following Smail et al. ( |
1998b). we parameterized subi source counts at 850 qn as: where the uncertainties are the lo errors. | 1998b), we parameterized submm source counts at 850 $\mu$ m as: where the uncertainties are the $\sigma$ errors. |
The slope a is. (1.1:0.2). aud thus. the average maeuification bias. (11.τὸν rauges ποια 0.9 to Ll. | The slope $\alpha$ is $-(1.1 \pm 0.2)$, and thus, the average magnification bias, $\langle B(f,z) \rangle$, ranges from 0.9 to 1.4. |
This estimate for the maeuification bias for 850 san sources with flux deusities in the rauge of 0.5 to 10 indy is slightly lower than what was previously considered (e... Blain 1997). | This estimate for the magnification bias for 850 $\mu$ m sources with flux densities in the range of 0.5 to 10 mJy is slightly lower than what was previously considered (e.g., Blain 1997). |
For the purpose of this calculation. where we are onlv interested in an upper huit to the redshift distribution. we apply the lowest possible amplification bias to all background sources. | For the purpose of this calculation, where we are only interested in an upper limit to the redshift distribution, we apply the lowest possible amplification bias to all background sources. |
This leads to an underestimated leusing rate aud an overestimated upper luit ou the background sourceredshift distribution. | This leads to an underestimated lensing rate and an overestimated upper limit on the background sourceredshift distribution. |
Since the redshitt distribution of subi sources. C(:). is unknown. we caleulate the observed umuuuber of lensed sources as a fuuction of (2). the effective average redshift uuder the assiuption that all sources are at this redshift: | Since the redshift distribution of submm sources, $C(z)$ , is unknown, we calculate the observed number of lensed sources as a function of $\langle z \rangle$ , the effective average redshift under the assumption that all sources are at this redshift: |
These early “Wari Spitzer” data were some of the first to be reduced with the new "Warm IRAC" pipeline. | These early “Warm Spitzer" data were some of the first to be reduced with the new “Warm IRAC" pipeline. |
The iuages were processed usiug the Spitzer Scieuce Center pipeline version 515.18.0. which does no inclucle tje artifact correction steps present in the cold mission pipelines. | The images were processed using the Spitzer Science Center pipeline version S18.18.0, which does not include the artifact correction steps present in the cold mission pipelines. |
Therefore. in addition o the normal Spitzer Science Center (55C) processing. the BCD (Basic Calibratec Data) frames were rtu through the Warm-Mission Cohunn PulldownC | Therefore, in addition to the normal Spitzer Science Center (SSC) processing, the BCD (Basic Calibrated Data) frames were run through the Warm–Mission Column Pulldown. |
orrector?.. (Àlakovoz&Wlan2005) was used with the staudard aud nane lists [or the overla2 Correction aid mosaicking. | \citep{2005ASPC..347...81M} was used with the standard and name lists for the overlap correction and mosaicking. |
The10st linportan objective ol this work is to p"OCuce al accurate calibration «X the Ceplieid period-lurMnosity relatiO1. | The most important objective of this work is to produce an accurate calibration of the Cepheid period–luminosity relation. |
Our shotometry is on he same scale as tle staucla«d calibration described jiu Reacretal.(2005).. | Our photometry is on the same scale as the standard calibration described in \citet{2005PASP..117..978R}. |
They define the [lux scae aud zero point such that tliey correspond to a I0-pixel radius apeture. with a sky aunuulus rurni18oO rou 12 to 20 pixels (iu 1nits of native IRAC pixels) wit1 pixel-phase and array-locatiou coctions applied as necessary. | They define the flux scale and zero point such that they correspond to a 10-pixel radius aperture, with a sky annulus running from 12 to 20 pixels (in units of native IRAC pixels) with pixel–phase and array–location corrections applied as necessary. |
Although the majority of the Cepheids in our sate are Isolated. several have €ose neighbors that fall inside tte 10-pixel aperture. | Although the majority of the Cepheids in our sample are isolated, several have close neighbors that fall inside the 10-pixel aperture. |
These stars cotld uot be measured correctly ising aperture photometry. | These stars could not be measured correctly using aperture photometry. |
lu order to reduce the entire LMC sample it a consistent way. Point Respouse Function (PREF) fitting was used in place of aperture photometry. utiliziug the SSC softw"ua'e 2005)... | In order to reduce the entire LMC sample in a consistent way, Point Response Function (PRF) fitting was used in place of aperture photometry, utilizing the SSC software \citep{2005PASP..117.1113M}. |
The point respouse ftnetion (PRE models tie detector respouse to a poiut source. and can be thought of as au over-saimpled representation of the point spread thnuctiou. | The point response function (PRF) models the detector response to a point source, and can be thought of as an over-sampled representation of the point spread function. |
The BCDs were mosaicked to obtain the bad jixel naps. bu the PREF fitting was perforined ou each BCD uxividually using the 1ale list. witl two changes. | The BCDs were mosaicked to obtain the bad pixel maps, but the PRF fitting was performed on each BCD individually using the name list, with two changes. |
First. rather thau using the PRE cor'espoudiug to the center of the arav. the set of 25 spatiallvy-dlepeudeut PREs was used [or eachuel*. | First, rather than using the PRF corresponding to the center of the array, the set of 25 spatially–dependent PRFs was used for each. |
. The IRAC detectors are kiowu to lave some spatial response variations: usiug he PRE ma2 eusured that any variation iu fhx was properly accounted for. | The IRAC detectors are known to have some spatial response variations; using the PRF map ensured that any variation in flux was properly accounted for. |
Second. the normalizatjon radius — the radius to which the fit is cak"ulatecd — was set to 1000 pixels. | Second, the normalization radius – the radius to which the fit is calculated – was set to 1000 pixels. |
As the PRE is sampled at 100 tiues per pixel. the 1000 pixel norijalizatiou radius corresponds to the calibration apertre radius. | As the PRF is sampled at 100 times per pixel, the 1000 pixel normalization radius corresponds to the calibration aperture radius. |
For each BCD. all the stars in eac1 inage were measured. | For each BCD, all the stars in each image were measured. |
The construction of the IRAC PRE is a complex task requiring a large number of observations ol a stable star. | The construction of the IRAC PRF is a complex task requiring a large number of observations of a stable star. |
Wari Spitzer PRE mocles are not yet available [rom the Spitzer Scieuce Center: however. it is still possible to obtain accurate photometry using the cold PRE. correcting as necessary lor the predicted differences beween the cold aud warum models. | Warm Spitzer PRF models are not yet available from the Spitzer Science Center; however, it is still possible to obtain accurate photometry using the cold PRF, correcting as necessary for the predicted differences between the cold and warm models. |
The mostsiguilicant | The mostsignificant |
where ly denotes the interstellar absorption in V. | where $A_V$ denotes the interstellar absorption in $V$. |
It has to be stressed that Chis can onlv be taken as a first estimate (see Table 6)). | It has to be stressed that this can only be taken as a first estimate (see Table \ref{fpmspulsinfo}) ). |
With the relations given in(1993).. the (5—yy colors of (he classical 9 Seuti stars were transformed into (4—V), values. | With the relations given in, the $(b-y)_0$ colors of the classical $\delta$ Scuti stars were transformed into $(B-V)_0$ values. |
Any. transformation is affected by errors. but they are smaller when transforming intermediate band to broad band photometry (han vice versa. | Any transformation is affected by errors, but they are smaller when transforming intermediate band to broad band photometry than vice versa. |
Consequently. we translormed the (5—η colors of the 0 Seuti stars to (D—VJo. | Consequently, we transformed the $(b-y)_0$ colors of the $\delta$ Scuti stars to $(B-V)_0$. |
It is now possible to compare classical 0. Scuti to pulsaüng PAIS stars in a common parameter space. Ady vs. (D—Vg. (Fig. 2)). | It is now possible to compare classical $\delta$ Scuti to pulsating PMS stars in a common parameter space, $M_V$ vs. $(B-V)_0$, (Fig. \ref{newstrip-BV}) ). |
For the 6 Sceuti stus My is taken from and for the PALS pulsators and candidates Aly was computed [rom the cluster distances. (filled ancl open diamonds: Table 5)) and parallaxes of the field stars (filled squares: Table 6)). | For the $\delta$ Scuti stars $M_V$ is taken from and for the PMS pulsators and candidates $M_V$ was computed from the cluster distances (filled and open diamonds; Table \ref{pmspulsinfo}) ) and parallaxes of the field stars (filled squares; Table \ref{fpmspulsinfo}) ). |
For the majority of PAIS field pulsators the published parallaxes either place the stars below the ZAMS or have large (1.6... > 1054) errors6). | For the majority of PMS field pulsators the published parallaxes either place the stars below the ZAMS or have large (i.e., $>$ ) errors. |
. Hence. only (wo pulsating PAIS field stars are selected for the IL. diagram: 2 Pie and ID 104237. with errors in the parallaxes of and6%... respectively. | Hence, only two pulsating PMS field stars are selected for the HR diagram: $\beta$ Pic and HD 104237, with errors in the parallaxes of and, respectively. |
According to2.. PAIS pulsators ancl o 5Scuti stars seem (o populate the same instability region in the HR diagram. | According to, PMS pulsators and $\delta$ Scuti stars seem to populate the same instability region in the HR diagram. |
A lack of pulsating PAIS stars is visible at the "eoo corner of the instability region for classical 6 Seuti stars. | A lack of pulsating PMS stars is visible at the “cool" corner of the instability region for classical $\delta$ Scuti stars. |
Whether this is onlv a selection effect caused by poor nuuiber statistics or has some as(voplvsical reason. can only be speculated about. | Whether this is only a selection effect caused by poor number statistics or has some astrophysical reason, can only be speculated about. |
Figure 2. also shows the uncertainties in My: and (2—Wj) lor the PAIS pulsators. | Figure \ref{newstrip-BV} also shows the uncertainties in $M_V$ and $(B-V)_0$ for the PMS pulsators. |
For the pulsating PAIScluster stars the errors in Ady were computed [rom the errors in distance (Table 5)) and the (2B—V), errors were either taken [rom the literature. if available. or propagated from the listed errors in V. and B measurements. | For the pulsating PMS stars the errors in $M_V$ were computed from the errors in distance (Table \ref{pmspulsinfo}) ) and the $(B-V)_0$ errors were either taken from the literature, if available, or propagated from the listed errors in $V$ and $B$ measurements. |
In an analogous manner. (he Ay errors for the pulsating PMSfield stars were caleulated [rom Che errors in the parallax (Table 6)). | In an analogous manner, the $M_V$ errors for the pulsating PMS stars were calculated from the errors in the parallax (Table \ref{fpmspulsinfo}) ). |
As the parallax errors ave twpically larger thanLOY... only two of the 15 pulsating PAIS Ποιά stars wilh accurate enough parallaxes could be used for the comparison of PAIS pulsators ancl 9 Seuli stars in the WR. diagram. | As the parallax errors are typically larger than, only two of the 18 pulsating PMS field stars with accurate enough parallaxes could be used for the comparison of PMS pulsators and $\delta$ Scuti stars in the HR diagram. |
For the pulsating PAIS field stars errors for (D—V) are scarce in the literature ancl we estimated them as standard. deviations of independently measured ancl published V. and D magnitudes. | For the pulsating PMS field stars errors for $(B-V)$ are scarce in the literature and we estimated them as standard deviations of independently measured and published $V$ and $B$ magnitudes. |
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