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19 may be approximated as The cooling coefficient 5 is expected to be dominated by svnchrotron and/or external Compton cooling. | \ref{Nfc} may be approximated as The cooling coefficient $\nu_0$ is expected to be dominated by synchrotron and/or external Compton cooling. |
With the magnetic-field scaling from Eq. I4. | With the magnetic-field scaling from Eq. \ref{B}, |
both wuβS4 0d d Carry a dependence x>I7. | both $u'_{\rm ext}$ and $u'_B$ carry a dependence $\propto
\Gamma^2$. |
LThus. we find. lor. the normalizationBH of. the ultrarelativistic. particle- population: In order to use Eq. | Thus, we find for the normalization of the ultrarelativistic particle population: In order to use Eq. |
22 for a light curve estimate. we assume. again. for simplicity. 6).0 and therefore Dz2T. | \ref{Fsy} for a light curve estimate, we assume, again, for simplicity, $\theta_{\rm obs} = 0$ and therefore $D \approx
2 \, \Gamma$. |
Consequently. the characteristic electron energy 5=ve'/egxET.f. | Consequently, the characteristic electron energy $\gamma = \sqrt{\epsilon' / \epsilon_B} \propto \Gamma^{-1}$. |
This vields an expected lightcurve decay in the fast-cooling svnchrotron regine of In particular. for an injection index of Β’=3.4. as inferred from (he optical spectral index. a light curve of Fy,xΒ£47 is expected. | This yields an expected lightcurve decay in the fast-cooling synchrotron regime of In particular, for an injection index of $q = 3.4$, as inferred from the optical spectral index, a light curve of $F_{\nu} \propto t_{\rm obs}^{-2.2}$ is expected. |
However. it should be pointed out that this can only be considered an upper limit to the steepness of the decay. | However, it should be pointed out that this can only be considered an upper limit to the steepness of the decay. |
Any non-zero observing angle will flatten the decay of the light curve as it introduces a shallower decay of the Doppler factor D wilh decreasing D and therefore with Gime. | Any non-zero observing angle will flatten the decay of the light curve as it introduces a shallower decay of the Doppler factor $D$ with decreasing $\Gamma$ and therefore with time. |
In any case. by the time the plasmoid is in (he deceleration phase. the monochromatic flux has already decreased by about an order of magnitude Irom its initial peak value. and is likely to be overwhelmed by other emission components in the jet. | In any case, by the time the plasmoid is in the self-similar deceleration phase, the monochromatic flux has already decreased by about an order of magnitude from its initial peak value, and is likely to be overwhelmed by other emission components in the jet. |
We do therefore not expect to observe the limiting deceleration case clirectly. | We do therefore not expect to observe the limiting deceleration case directly. |
We only developed (this analvtical case to demonstrate (he agreement of our nunmerical simulations with the analvtical expectation in the following section. | We only developed this analytical case to demonstrate the agreement of our numerical simulations with the analytical expectation in the following section. |
In order to hishlisht the salient features of our model. we assume a simple cvlindrical jet geometry. with a constant cross section of the jet. =zA. as well as a homogeneous external medium with density p(r)=pea. | In order to highlight the salient features of our model, we assume a simple cylindrical jet geometry with a constant cross section of the jet, $A \equiv
\pi R_b^2$, as well as a homogeneous external medium with density $\rho (r)
\equiv \rho_{\rm ext}$. |
We achieved good fits to the observed optical hight curves and overall SED shape of 3C 279 with the parameters listed in Table 1.. | We achieved good fits to the observed optical light curves and overall SED shape of 3C 279 with the parameters listed in Table \ref{parameters}. . |
Our choice of the external density corresponds to a number density of meg=100 7. | Our choice of the external density corresponds to a number density of $n_{\rm ext} = 100$ $^{-3}$ . |
Thisis | Thisis |
the atmosphere in fairly good abundance. | the atmosphere in fairly good abundance. |
The condensation curve presented in C03 shows that the base of gehlenite is situated much deeper in the atinosphere than that of forsterite vielding into a larger verticle size of the cloud. | The condensation curve presented in C03 shows that the base of gehlenite is situated much deeper in the atmosphere than that of forsterite yielding into a larger verticle size of the cloud. |
Further. iron with higher refractivitv may undercircle most of the silicate elouds. | Further, iron with higher refractivity may undercircle most of the silicate clouds. |
Inclusion of all these will lead to substantial increase in degree of polarization and therefore. much smaller grain size may be needed in order (o explain the observed degree of polarization. | Inclusion of all these will lead to substantial increase in degree of polarization and therefore, much smaller grain size may be needed in order to explain the observed degree of polarization. |
Lastly. multiple scattering will lead to much less polarization ancl therefore in order to lit the observed polarization one has to either consider much higher oblateness of the objects or has (o increase the grain number density substantially. | Lastly, multiple scattering will lead to much less polarization and therefore in order to fit the observed polarization one has to either consider much higher oblateness of the objects or has to increase the grain number density substantially. |
An increase in grain number density needs smaller grain size which may contraclict the present theoretical understanding on the nature and formation of dust grain. | An increase in grain number density needs smaller grain size which may contradict the present theoretical understanding on the nature and formation of dust grain. |
On the other hand. differential photometric observation of several L cwarls could not detect any non-periodic variability of many objects that show hieh polarization. | On the other hand, differential photometric observation of several L dwarfs could not detect any non-periodic variability of many objects that show high polarization. |
This indicates optically thin dust shell and therefore polarization by single dust scattering is equite reasonable. | This indicates optically thin dust shell and therefore polarization by single dust scattering is quite reasonable. |
We have investigated the optical linear polarization from L dwarls of different spectral (vpe by considering single dust scattering. | We have investigated the optical linear polarization from L dwarfs of different spectral type by considering single dust scattering. |
Forsterite is considered to be the dominant species among (he various condensates (hat could be present in the atmosphere of L dwarls. | Forsterite is considered to be the dominant species among the various condensates that could be present in the atmosphere of L dwarfs. |
The location of the cloud base ancl Che cloud deck is determined from (he condensation curve for forsterite aud (he atmospheric temperature-pressure profiles of dillerent spectral types. | The location of the cloud base and the cloud deck is determined from the condensation curve for forsterite and the atmospheric temperature-pressure profiles of different spectral types. |
The surface eravily is fixed at 10? em 7 and a wide range of rotational velocity is considered. | The surface gravity is fixed at $10^5$ cm $^{-2}$ and a wide range of rotational velocity is considered. |
It is found that the degree of linear polarization decreases from hotter to cooler L clwarls. | It is found that the degree of linear polarization decreases from hotter to cooler L dwarfs. |
llowever. L dwarls with effective temperature greater than 2200 Ix should not show detectable amount of polarization due to dust scattering as most of the dust would either evaporate from the atmosphere or condensation is not favored at such high temperature. | However, L dwarfs with effective temperature greater than 2200 K should not show detectable amount of polarization due to dust scattering as most of the dust would either evaporate from the atmosphere or condensation is not favored at such high temperature. |
It is found that the mean diameter of grains that is consistent with the observed polarization should not exceed a few micron although a small amount of very. large grains at the base of the eloud for comparatively cooler L dsvarfs may well be accommodated. | It is found that the mean diameter of grains that is consistent with the observed polarization should not exceed a few micron although a small amount of very large grains at the base of the cloud for comparatively cooler L dwarfs may well be accommodated. |
Lowever. the observational data of 2ATASS J00364-1821 clearly indicates the presence of sub-micron size grain. | However, the observational data of 2MASS J0036+1821 clearly indicates the presence of sub-micron size grain. |
Further polarimetric observation at (he optical and al other wavelengths would provide convincing inlormation on the properties and distribution of dust in the atmosphere of L dwarts. | Further polarimetric observation at the optical and at other wavelengths would provide convincing information on the properties and distribution of dust in the atmosphere of L dwarfs. |
We are thankful to the referee [or several valuable comments and suggestions. | We are thankful to the referee for several valuable comments and suggestions. |
The presence of ionized material in the interstellar medium (hereafter. 15M) can be attributed ΞΏ two distinctive mechanisms. | The presence of ionized material in the interstellar medium (hereafter, ISM) can be attributed to two distinctive mechanisms. |
List. photoionization of the. surrouncling neutral gas by the strong. energetic ultra-violet (UV) Dux of nearby massive stars is largeh βresponsible for the detection of the standard. hydrogen. Bamer series. twpically usec for the morphological and. kinenivical description of regions. | First, photoionization of the surrounding neutral gas by the strong, energetic ultra-violet (UV) flux of nearby massive stars is largely responsible for the detection of the standard hydrogen Balmer series, typically used for the morphological and kinematical description of regions. |
Secondly. stelar winds with high terminal velocities ancl violent Supernova blasts are. commonly associated with the propagaticn of transonic and supersonic shock waves in the surrouning medium. | Secondly, stellar winds with high terminal velocities and violent supernova blasts are commonly associated with the propagation of transonic and supersonic shock waves in the surrounding medium. |
The important increase of the. post-shocked eas temperature favors its ionization through shock exciation. | The important increase of the post-shocked gas' temperature favors its ionization through shock excitation. |
Pionecring work bw SabbacinΞΏαΌ±af.(1977) has compared specific Lux ratios fXx a variety of ionic transitions in the ISAL optical gas. | Pioneering work by \citet{Sab1977} has compared specific flux ratios for a variety of ionic transitions in the ISM optical gas. |
This allowed the authors to approximately separate stancdare regions ancl planetary. nebulae (hereafter. PNe) mostly governed by photoionization from shock-dominated supernova remnants (hereafter. SNRs). | This allowed the authors to approximately separate standard regions and planetary nebulae (hereafter, PNe) mostly governed by photoionization from shock-dominated supernova remnants (hereafter, SNRs). |
These diagnostic diagrams were often used. in literature. to classifv large amount of ionized targets in large-scale objects. for example more-or-less distant galaxies (c.g... MagriniΞΏαΌ±ad.2003:lΓΌesgo&LΓ³pez 2006)). | These diagnostic diagrams were often used, in literature, to classify large amount of ionized targets in large-scale objects, for example more-or-less distant galaxies (e.g., \citealt{Mag2003,Rie2006}) ). |
Fhis has led to emission-line ratio plots in which each object is usually statistically represented bv a single point. | This has led to emission-line ratio plots in which each object is usually statistically represented by a single point. |
Obviously. intrinsic variations. within a given object. can be investigated. by. targeting Galactic objects individually (e.g. Phillips&Cuesta1998.1999:PhillipsefaΒ£, 2010)). | Obviously, intrinsic variations, within a given object, can be investigated by targeting Galactic objects individually (e.g., \citealt{Phi1998,Phi1999,Phi2010}) ). |
This allows to spatially resolve much smaller (Β«00.1 pe) structures ancl artifacts characterized by peculiar line ratios that are. otherwise. unrevealed or statistically negligible in observations using poorer angular resolutions. | This allows to spatially resolve much smaller $\ll$ 0.1 pc) structures and artifacts characterized by peculiar line ratios that are, otherwise, unrevealed or statistically negligible in observations using poorer angular resolutions. |
The investigation of close ionized objects has alreacy revealed that photoionization and shock excitation can both be found in individual regions. | The investigation of close ionized objects has already revealed that photoionization and shock excitation can both be found in individual regions. |
Ixinematically speaking. standard | Kinematically speaking, standard |
>3M. Stefanoetal.2004:Williams2004). | $>$$_{\odot}$ \citealp{mcclintock2004}, \citep{williams2004}. \citep{kong2002,distefano2004,williams2004}. |
. 7/5T 1257 | $HST$ \citep{williams2004,williams2005bh1,williams2005bh2}. |
2=25.5. | $HST$ $B=25.5$ |
Valerie Myerscough Prize. | Valerie Myerscough Prize. |
DC and LC acknowledge the support of the Australian Research Council. | BC and LC acknowledge the support of the Australian Research Council. |
(he midplane) must be cool and dense. whereas the bottom lid the one closest to the midplane) must be warm and lisht. | the midplane) must be cool and dense, whereas the bottom lid the one closest to the midplane) must be warm and light. |
These thermal lids will be weakened or destroved either bv radiation of heat or turbulent erosion. | These thermal lids will be weakened or destroyed either by radiation of heat or turbulent erosion. |
Thus. a vertical [low must develop within the vortex (o maintain (he lids. | Thus, a vertical flow must develop within the vortex to maintain the lids. |
For a high-pressure antievelone centered on the midplane. gas will rise away [rom the midplane. adiabatically expanding and cooling. | For a high-pressure anticyclone centered on the midplane, gas will rise away from the midplane, adiabatically expanding and cooling. |
At the lids. the gas diverges outward and recirculates back toward (he midplane along the edges of (he vortex. | At the lids, the gas diverges outward and recirculates back toward the midplane along the edges of the vortex. |
For an antΓΌevclone that is completely above the midplane. gas will rise (sink) toward (he top (bottom) lid. where it will expand and cool (compress ancl heat-up). | For an anticyclone that is completely above the midplane, gas will rise (sink) toward the top (bottom) lid, where it will expand and cool (compress and heat-up). |
We can use (he arguments in (his section to derive simple scaling relationships for the vortex flow variables. which will then guide us in making further approximations to the {hud equations. | We can use the arguments in this section to derive simple scaling relationships for the vortex flow variables, which will then guide us in making further approximations to the fluid equations. |
We decompose each flow variable into a ΓΌme-independent. axisvimnietric component describing the base disk flow. which we will denote with overbars. and a component describing the vortex [low. which we will denote with tildes p=ptp.p=pt+p). | We decompose each flow variable into a time-independent, axisymmetric component describing the base disk flow, which we will denote with overbars, and a time-dependent component describing the vortex flow, which we will denote with tildes $\rho = \bar{\rho}+\tilde{\rho}$, $p = \bar{p}+\tilde{p}$ ). |
As discussed in the previous subsection. we assume that the horizontal vortex [low is subsonic: α½ΞΏ~e<1. | As discussed in the previous subsection, we assume that the horizontal vortex flow is subsonic: $\tilde{v}_{\bot}/c_s\sim\epsilon<1$. |
IE the Rossby number is of order unity or less. then the horizontal pressure force must be of the same order as the Coriolis force (geostrophic balance): p/pA,~206. where pis the pressure perturbation associated with the vortex. and pis the mean density. | If the Rossby number is of order unity or less, then the horizontal pressure force must be of the same order as the Coriolis force (geostrophic balance): $\tilde{p}/\bar{\rho}\Lambda_r\sim 2\Omega\tilde{v}_{\bot}$, where $\tilde{p}$ is the pressure perturbation associated with the vortex, and $\bar{\rho}$ is the mean density. |
Using and(2-6).. one can show that the pressure perturbation in the vortex must scale as: p/pc�/Ho. | Using and, one can show that the pressure perturbation in the vortex must scale as: $\tilde{p}/\bar{p} \sim \epsilon^2/Ro$. |
We also assume the vertical buovancy force is of order the vertical pressure gradient: q.T/T.~p/pll. which leads to the temperature fluctuation having the same scaling as the pressure [lnctuation: T/T~p/p. | We also assume the vertical buoyancy force is of order the vertical pressure gradient: $g_z\tilde{T}/\bar{T}\sim\tilde{p}/\bar{\rho}H$, which leads to the temperature fluctuation having the same scaling as the pressure fluctuation: $\tilde{T}/\bar{T}\sim\tilde{p}/\bar{p}$. |
Because of the ideal eas law. this implies that. density fluctuations also have the same scaling. | Because of the ideal gas law, this implies that density fluctuations also have the same scaling. |
The scaling for the vertical velocity is set bv. the maintenance of temperature perturbations in (he thermal lids. | The scaling for the vertical velocity is set by the maintenance of temperature perturbations in the thermal lids. |
In the temperature equation. temperature fluctuations are created via pressure-volume work: P.T(ds/dz)/C.. where 6. is the vertical velocitv. $is the mean entropy profile. and. C. is the heat capacity at constant volume. | In the temperature equation, temperature fluctuations are created via pressure-volume work: $\tilde{v}_z\bar{T}(d\bar{s}/dz)/C_v$, where $\tilde{v}_z$ is the vertical velocity, $\bar{s}$ is the mean entropy profile, and $C_v$ is the heat capacity at constant volume. |
The destruction process is either radiative heating/cooling. T/Teoot: where To 1s a characteristic cooling Gime: or turbulent advection. T/s,ddyviT/A. where 7,44 15 a characteristic eddy diffusion time. | The destruction process is either radiative heating/cooling, $\tilde{T}/\tau_{cool}$, where $\tau_{cool}$ is a characteristic cooling time; or turbulent advection, $\tilde{T}/\tau_{eddy}\sim\tilde{v}_{\bot}\tilde{T}/\Lambda_r$, where $\tau_{eddy}$ is a characteristic eddy diffusion time. |
If the vortex flow is smooth ancl lznninar. (hen we balance pressure-volume work with thermal cooling to obtain the scaling: 6.~(β¬/Ro)H7,4. | If the vortex flow is smooth and laminar, then we balance pressure-volume work with thermal cooling to obtain the scaling: $\tilde{v}_z\sim (\epsilon^2/Ro)H/\tau_{cool}$. |
Thus. if the vortex is laminar and the cooling time is long. the vertical velocity can be quite small. | Thus, if the vortex is laminar and the cooling time is long, the vertical velocity can be quite small. |
On (he other hand. if the vortex flow is highly turbulent. we balance pressure-volume work ancl the (turbulent advection of heat to obtain β ~ D>: : βββ²βββ½βΊβ’ββ§β΄βββΈββΆβ£β½βΌββ΄βΏβ΄β€ββ
βΊββΌβ²βββ½βΌβ²βββ½βΊβ’ββ§β΄βββΈβ½ββ΄ββ
βΌβ²β₯ββ§β΄ββͺββββ½βββ½βββ½ββ―β²βββ½ββββββ§β΄ββ
ββ βΌβ²βΊββββΊβββ΄β£β½β»β₯βΌβ²ββ»β
β
: : | On the other hand, if the vortex flow is highly turbulent, we balance pressure-volume work and the turbulent advection of heat to obtain the scaling: $\tilde{v}_z\sim\epsilon^2$ .These scaling relationships are summarized in Table \ref{T:scalings}. . |
On (he other hand. if the vortex flow is highly turbulent. we balance pressure-volume work ancl the (turbulent advection of heat to obtain β ~ D>: : βββ²βββ½βΊβ’ββ§β΄βββΈββΆβ£β½βΌββ΄βΏβ΄β€ββ
βΊββΌβ²βββ½βΌβ²βββ½βΊβ’ββ§β΄βββΈβ½ββ΄ββ
βΌβ²β₯ββ§β΄ββͺββββ½βββ½βββ½ββ―β²βββ½ββββββ§β΄ββ
ββ βΌβ²βΊββββΊβββ΄β£β½β»β₯βΌβ²ββ»β
β
: : | On the other hand, if the vortex flow is highly turbulent, we balance pressure-volume work and the turbulent advection of heat to obtain the scaling: $\tilde{v}_z\sim\epsilon^2$ .These scaling relationships are summarized in Table \ref{T:scalings}. . |
On (he other hand. if the vortex flow is highly turbulent. we balance pressure-volume work ancl the (turbulent advection of heat to obtain β ~ D>: : βββ²βββ½βΊβ’ββ§β΄βββΈββΆβ£β½βΌββ΄βΏβ΄β€ββ
βΊββΌβ²βββ½βΌβ²βββ½βΊβ’ββ§β΄βββΈβ½ββ΄ββ
βΌβ²β₯ββ§β΄ββͺββββ½βββ½βββ½ββ―β²βββ½ββββββ§β΄ββ
ββ βΌβ²βΊββββΊβββ΄β£β½β»β₯βΌβ²ββ»β
β
: : : | On the other hand, if the vortex flow is highly turbulent, we balance pressure-volume work and the turbulent advection of heat to obtain the scaling: $\tilde{v}_z\sim\epsilon^2$ .These scaling relationships are summarized in Table \ref{T:scalings}. . |
On (he other hand. if the vortex flow is highly turbulent. we balance pressure-volume work ancl the (turbulent advection of heat to obtain β ~ D>: : βββ²βββ½βΊβ’ββ§β΄βββΈββΆβ£β½βΌββ΄βΏβ΄β€ββ
βΊββΌβ²βββ½βΌβ²βββ½βΊβ’ββ§β΄βββΈβ½ββ΄ββ
βΌβ²β₯ββ§β΄ββͺββββ½βββ½βββ½ββ―β²βββ½ββββββ§β΄ββ
ββ βΌβ²βΊββββΊβββ΄β£β½β»β₯βΌβ²ββ»β
β
: : : | On the other hand, if the vortex flow is highly turbulent, we balance pressure-volume work and the turbulent advection of heat to obtain the scaling: $\tilde{v}_z\sim\epsilon^2$ .These scaling relationships are summarized in Table \ref{T:scalings}. . |
On (he other hand. if the vortex flow is highly turbulent. we balance pressure-volume work ancl the (turbulent advection of heat to obtain β ~ D>: : βββ²βββ½βΊβ’ββ§β΄βββΈββΆβ£β½βΌββ΄βΏβ΄β€ββ
βΊββΌβ²βββ½βΌβ²βββ½βΊβ’ββ§β΄βββΈβ½ββ΄ββ
βΌβ²β₯ββ§β΄ββͺββββ½βββ½βββ½ββ―β²βββ½ββββββ§β΄ββ
ββ βΌβ²βΊββββΊβββ΄β£β½β»β₯βΌβ²ββ»β
β
: : : [ | On the other hand, if the vortex flow is highly turbulent, we balance pressure-volume work and the turbulent advection of heat to obtain the scaling: $\tilde{v}_z\sim\epsilon^2$ .These scaling relationships are summarized in Table \ref{T:scalings}. . |
On (he other hand. if the vortex flow is highly turbulent. we balance pressure-volume work ancl the (turbulent advection of heat to obtain β ~ D>: : βββ²βββ½βΊβ’ββ§β΄βββΈββΆβ£β½βΌββ΄βΏβ΄β€ββ
βΊββΌβ²βββ½βΌβ²βββ½βΊβ’ββ§β΄βββΈβ½ββ΄ββ
βΌβ²β₯ββ§β΄ββͺββββ½βββ½βββ½ββ―β²βββ½ββββββ§β΄ββ
ββ βΌβ²βΊββββΊβββ΄β£β½β»β₯βΌβ²ββ»β
β
: : : [ | On the other hand, if the vortex flow is highly turbulent, we balance pressure-volume work and the turbulent advection of heat to obtain the scaling: $\tilde{v}_z\sim\epsilon^2$ .These scaling relationships are summarized in Table \ref{T:scalings}. . |
On (he other hand. if the vortex flow is highly turbulent. we balance pressure-volume work ancl the (turbulent advection of heat to obtain β ~ D>: : βββ²βββ½βΊβ’ββ§β΄βββΈββΆβ£β½βΌββ΄βΏβ΄β€ββ
βΊββΌβ²βββ½βΌβ²βββ½βΊβ’ββ§β΄βββΈβ½ββ΄ββ
βΌβ²β₯ββ§β΄ββͺββββ½βββ½βββ½ββ―β²βββ½ββββββ§β΄ββ
ββ βΌβ²βΊββββΊβββ΄β£β½β»β₯βΌβ²ββ»β
β
: : : [ β
| On the other hand, if the vortex flow is highly turbulent, we balance pressure-volume work and the turbulent advection of heat to obtain the scaling: $\tilde{v}_z\sim\epsilon^2$ .These scaling relationships are summarized in Table \ref{T:scalings}. . |
We performed a series of tests for. potential svsteniatic photometric errors that may. alfect our. clustering results. | We performed a series of tests for potential systematic photometric errors that may affect our clustering results. |
Following the discussion in Blake et ((2007). we compared the angular correlation function measured for the "default sample with that obtained by restricting or extending the ealaxy selection in the following wavs: We refer the reader to Blake et ((2007) for a more thorough discussion of the possible effects. of these systematic errors. | Following the discussion in Blake et (2007), we compared the angular correlation function measured for the βdefaultβ sample with that obtained by restricting or extending the galaxy selection in the following ways: We refer the reader to Blake et (2007) for a more thorough discussion of the possible effects of these systematic errors. |
Our results are presented in Figures S.. 9. and 10.. | Our results are presented in Figures \ref{figsys1}, , \ref{figsys2} and \ref{figsys3}. |
Each plot is composed of four panels. one for each redshift slice. | Each plot is composed of four panels, one for each redshift slice. |
In each panel we show the angular correlation function for the default sample together with that corresponding to Γ change in the galaxy selection criteria. | In each panel we show the angular correlation function for the default sample together with that corresponding to a change in the galaxy selection criteria. |
We divide all the correlation functions by a power-law fit to the default measurement to render the results more easily comparable. | We divide all the correlation functions by a power-law fit to the default measurement to render the results more easily comparable. |
We conclude from Figures Ξ΄ and 9 that our results are robust against the details of the angular selection function: varvingIn dust extinction. seeing.α½Ξ½ overlap regionso and bright object masks all have little elfect on the measured correlation function. | We conclude from Figures \ref{figsys1} and \ref{figsys2} that our results are robust against the details of the angular selection function: varying dust extinction, seeing, overlap regions and bright object masks all have little effect on the measured correlation function. |
Figure LO reveals that the details of the star-galaxy separation allect the amplitude of the measured correlation function although not (to first order) the shape. | Figure \ref{figsys3} reveals that the details of the star-galaxy separation affect the amplitude of the measured correlation function although not (to first order) the shape. |
This amplitude shift is already. encoded in the stellar contamination [actor (1[Y | This amplitude shift is already encoded in the stellar contamination factor $(1-f)^2$. |
ln no case does a change in the ealaxy selection alter the cletectability or shape of the halo mocel signature. | In no case does a change in the galaxy selection alter the detectability or shape of the halo model signature. |
We have measured the angular correlation function of LRGs in the SDSS imaging survey. using accurate. photomoetric redshifts to divide the galaxies into narrow redshift. slices and create volume-lIimitedsamples. | We have measured the angular correlation function of LRGs in the SDSS imaging survey, using accurate photometric redshifts to divide the galaxies into narrow redshift slices and create volume-limitedsamples. |
We find that: | We find that: |
subnun position is not negligible: c.g. there would be ~4 A- sources (AgiΒ«23) on average within ar Β«S radius (see below). | submm position is not negligible; e.g. there would be $\sim 4$ $K$ -band sources $K_\mathrm{AB} < 23)$ on average within a $r<8''$ radius (see below). |
However. if SMGs have rather confined optical-NIIS colours and the surface density of objects with similar colours is reasonably low. it will be possible to. directly identifv SAIGs without using costly high-resolution submunm images. | However, if SMGs have rather confined optical-NIR colours and the surface density of objects with similar colours is reasonably low, it will be possible to directly identify SMGs without using costly high-resolution submm images. |
As in the radio identification method (Downesetal. 1986).. we can calculate the formal significance. of cach optical/NIB. association. given the number counts of colour-selected galaxies. | As in the radio identification method \citep{1986MNRAS.218...31D}, we can calculate the formal significance of each optical/NIR association, given the number counts of colour-selected galaxies. |
To do this we need a firm basis for the colours of SAIGs. | To do this we need a firm basis for the colours of SMGs. |
Since the discovery. of SALGs. extremely ΟΞΏαΏ¦ objects (EROs) usually defined. with (4AQvoun or (RA)vounZ5 and distant red galaxies with (JzdAvo2e)J rave been paid special attention as candidates of SAIC1 counterparts (e.g.Smaletal.1999.2002:Webb2003:Popeetal. 2005).. | Since the discovery of SMGs, extremely red objects (EROs) usually defined with $(I-K)_\mathrm{Vega}\ga 4$ or $(R-K)_\mathrm{Vega}\ga 5$ and distant red galaxies with $(J-K)_\mathrm{Vega}>2.3$ have been paid special attention as candidates of SMG counterparts \citep[e.g.][]{1999MNRAS.308.1061S,2002MNRAS.331..495S,2003ApJ...597..680W,2004MNRAS.354..193C,2004AJ....127..728F,2004MNRAS.351..447C,2005MNRAS.358..149P}. |
Ht turns out that the fraction of EROs or DRGs in SMCGs is not very high and their surface densities are not low enough to reject random associations. | It turns out that the fraction of EROs or DRGs in SMGs is not very high and their surface densities are not low enough to reject random associations. |
Among he NiR-selected galaxy populations in the literature. DzlIx-selected. star-forming galaxies (DzlxsDaclelietal.2004) may be the most promising counterparts of SALGs (Iteddyetal.2005:Bertoldi2007:Takagi 2007).. which lie at 1.432:22:2.5- and include heavily obscured galaxies. such as SALGs. as an extreme subset. | Among the NIR-selected galaxy populations in the literature, BzK-selected star-forming galaxies \citep[BzKs --][]{2004ApJ...617..746D} may be the most promising counterparts of SMGs \citep{2005ApJ...633..748R,2007ApJS..172..132B,2007MNRAS.381.1154T}, which lie at $1.4\la z \la 2.5$ and include heavily obscured galaxies, such as SMGs, as an extreme subset. |
In this paper. we indeed find that almost all of the A- (indicating higher redshifts) radio-detected SMCGs are i:dxs. | In this paper, we indeed find that almost all of the $K$ -faint (indicating higher redshifts) radio-detected SMGs are BzKs. |
We then use Bzlxs as a ΞΞΏΞ½ galaxy population to identify radio-uncdetectod SALGs at 144527:2.5. | We then use BzKs as a key galaxy population to identify radio-undetected SMGs at $1.4\la z \la 2.5$. |
Since the 3MIx-selection is based on observed wavelengths covering a spectral break atAA. we can naturally extend. this selection. technique to the higher redshift range. e.g. 2Z93. in the future by using a combination of wavebands at longer. wavelengths. | Since the BzK-selection is based on observed wavelengths covering a spectral break at, we can naturally extend this selection technique to the higher redshift range, e.g. $z\gg 3$, in the future by using a combination of wavebands at longer wavelengths. |
We also emphasize that the understanding of the physical relation between SALGs and. Dzlxs would be important. to reveal the evolution of galaxies at z~2. given that they have similar stellar masses of ~1045 MM. and. spatial correlation lengths of ~ 10h+ MMpe (Blainetal.2004:lxongetal.2006:Lavashict 2007).. | We also emphasize that the understanding of the physical relation between SMGs and BzKs would be important to reveal the evolution of galaxies at $z\sim 2$, given that they have similar stellar masses of $\sim 10^{11}$ $_\odot$, and spatial correlation lengths of $\sim 10$ $h^{-1}$ Mpc \citep{2004ApJ...611..725B,2006ApJ...638...72K,2007ApJ...660...72H}. |
We discuss. the hypothesis of SMGs being mereing Dzlxs. | We discuss the hypothesis of SMGs being merging BzKs. |
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