LWS research:Understanding the sources of the solar spectral and total irradiance variability and forecasting tools

2007/12/11

PI: J. Fontenla

LASP – Univ. of Colorado

Juan.Fontenla@lasp.colorado.edu

SRPM Project Goals•Diagnosis of physical conditions through the solar atmosphere; energy balance of radiative losses and mechanical heating.

•Evaluating proposed physical processes to determine the solar atmosphere structure and spectrum at all spatial and temporal scales.

•Synthesizing solar irradiance spectrum and its variations to improve the above and produce complete and quantitative physical models.

•Forecasting spectral irradiance at any time and position in the Heliosphere. Weekly and monthly forecast is now becoming possible.

SRPM Flow Scheme

Emitted Spectrum

Physical Models& Processes

ObservedSpectrum

Intermediate Parameters I(λ,μ,φ,t)

T,ne,nh,U,...(x,y,z,t)

nlev,nion,…(x,y,z,t)

I(λ,μ,φ,t)

0.8

0.6

0.4

0.2

0.0

Ion

izat

ion

Fra

ctio

n

104

2 3 4 5 6 7 8 9

105

2 3 4 5 6 7 8 9

106

Temperature (K)

Carbon Ionization and Mass FLow......... Static case (w/dif)_____ Upflow case (w/dif)

SRPM Technology• Full non-LTE radiative transfer for all relavant

species (including optically thick and thin)• Multi-dimensional radiative transfer, 1D and 3D• Modular, client-server, distributed structure• Extensive relational SQL database storage for:

– Atomic and molecular data

– Physical models and simulations

– Intermediate data (e.g., level populations)

• Object Oriented C++ reusable production code• I/O interfaces to text, binary, FITS, NETCDF• Parallel computing using available libraries

Modeling for various plasma regimes• Photosphere (using average 1D models and external 3D simulations)

– Slow motions (few km/s) dominated by convection overshoot – Weak ionization– All particles are unmagnetized– Plasma beta > 1– At or near LTE

• Chromosphere (using average 1D models and 3D MHD simulations)– Motions and inhomogeneities change from weak to strong– Weak ionization (np<<en~10-4 nH)– Ions unmagnetized, electrons magnetized (implies tensor conductivity)– Plasma beta crosses 1 somewhere within the chromosphere– Needs to consider full non-LTE radiative transfer radiative losses

• Corona (will use results from groups carrying coronal loops modeling)– Motions and ihomogeneities are very strong– Highly ionized– All species are magnetized– Plasma beta << 1– Non-LTE effects are extreme and but optically thin applies– Particle transport is large and probably important departures from Maxwellian

Boundary conditions between layers• Between photosphere and upper chromosphere:

– The low chromosphere is near radiative equilibrium– Driven by convective overshoot and also by Lorentz forces (i.e.

magnetic fields) in some locations– NLTE effects driven by illumination from above and below.

• Between corona and chromosphere:– The transition-region behaves like a boundary layer at the footpoints of

coronal loops or solar wind open field lines– Energy balance between energy carried by conduction and diffusion

from the corona is dissipated by radiation in the transition-region, optically thick and thin depending on species

– Mass also flows through the transition-region and supplies the solar wind

• (Cool loops exist embedded in the corona and are dynamic, e.g. spicules, but are not too important for the solar irradiance)

• (Warm loops exist embedded in the chromosphere and are dynamic, but are not too important for the solar irradiance)

Photosphere (radiation/convection)

500 nm 800 nm 1200 nm 1600 nm

Stein & Nordlund 2000 convection simulations snapshotsSRPM absolute radiance, wavelength and CLV dependence

Slit spectrum

1 103

1 104

1 105

1 106

0

200

400

SRPM 306Stein & NordlundSRPM 306 + 30 km

Pressure (dyne cm^-2)

Hei

ght (

km)

Comparison of spatial averages with semi-empirical modelspoints to improvements in average models and in simulations

Mg I 4572C I 5381 CN band

1 103

1 104

1 105

1 106

4000

6000

8000

SRP M 306Stein & NordlundSRP M 306 * 0.95

Pressure (dyne cm^-2)

Tem

pera

ture

(K)

Solar Chromosphere (radiation/plasma heating?)

7000

6000

5000

4000

3000

Tem

per

atu

re (

K )

SRPM 305 C1 VAL B COmosphere

3000

2500

2000

1500

1000

500

0

-500

Hei

ght

( k

m )

10-2

10-1

100

101

102

103

104

105

106

Gas Pressure (dyne cm-2)

SRPM 305 C1 VAL B COmosphere

5500

5000

4500

4000

Bri

ghtn

ess

Tem

pera

ture

( K

)

4.467 04.46684.46664.46644.46624.4660

Wavelength ( mm )

Farmer & Norton

5100

5000

4900

4800

4700

4600

4500

4400

Bri

ghtn

ess

Tem

per

atu

re (

K )

150014801460144014201400

Wavelength ( Å )

SRPM 305 Curdt et al.

8000

7000

6000

5000

4000

Bri

ghtn

ess

Tem

per

atu

re (

K )

0.012 4 6 8

0.12 4 6 8

12 4 6 8

10

Wavelength ( mm )

SRPM 305 Urpo et al. (1987) Loukitcheva et al. (2004) Beckman et al. (1973) Degiacomi & Kneubuehl (1985) Fürst (1980) Boreiko & Clark (1987)

New intranetwork model (B) matches theobservations at most λ with no bifurcation. Allows a simple average model for computing all wavelengths.

Comparison of semi-empirical quiet-Sun model spectrum with observations, shows a good

match but also some details to improve

300 320 340 360 380 400 420 440 460 480 500 520 540 560 580 6005200

5400

5600

5800

6000

6200

SIMSOLSPECSRPM 306

Wavelength (nm)

Bri

ghtn

ess

Tem

pera

ture

(K

)

600 800 1000 1200 1400 1600 1800 2000 2200 2400 2600 2800 30005600

5800

6000

6200

6400

6600

SIMSOLSPECSRPM 306

Wavelength (nm)

Bri

ghtn

ess

Tem

pera

ture

(K

)

5880 5882 5884 5886 5888 5890 5892 5894 5896 5898 5900 59020

1 106

2 106

3 106

4 106

SRPM 306KP atlas

Wavelength (A)

Inte

nsi

ty

4572 4572.5 4573 4573.5 45740

1 106

2 106

3 106

4 106

5 106

SRPM 306KP atlas

Wavelength (A)

Inte

nsi

ty

6560 6565 6570 65750

1 106

2 106

3 106

SRPM 306KP atlas

Wavelength (A)

Inte

nsi

ty

4302 4304 4306 4308 4310 43120

2 106

4 106

6 106

SRPM 306KP atlas

Wavelength (A)

Inte

nsi

ty

3880 3881 3882 3883 3884 38850

2 106

4 106

6 106

SRPM 306KP atlas

Wavelength (A)

Inte

nsi

ty

4.37 106

1.486 105

irr1

irrna

38853880 w1 wna

4.464 4.466 4.468 4.47

6000

8000

SRPM 306Farmer & Norton

Wavelength (micron)

Inte

nsi

ty

16.01 16.015 16.02 16.02550

52

54

56

58

60

SRPM 306Farmer & Norton

Wavelength (micron)

Inte

nsi

ty

4855 4860 48650

1 106

2 106

3 106

4 106

5 106

SRPM 306KP atlas

Wavelength (A)

Inte

nsi

ty

H alpha Na I D lines H beta Mg I 4572 & Ti II 4573

CN Band head CH Band (G-band) CO Bands OH Lines

Model 305 spectrum is ~3% too bright compared with the current observations of spectral irradiance. but the observations error is comparable.

Upper chromospheric network intensity structure shows

distributionwith relationship to magnetic fields

UV (1540 A) continuum MDI magnetogram

The network intensity distributionis log-normal, an additional tail appears in active

regions, we model a discretized distribution

UV continuumLyα

Ca II K3 Red cont.

Chromospheric heating & “microturbulence” appear to be closely related

Model 305-306 gives:Lower chromosphere: decreasing T - radiative equilibrium – subsonic motions -Vturb 1-3 km/sUpper chromosphere: relatively high T plateau - strong UV losses and heating – near-sonic motions - Vturb > 9 km/s

Heavy ions dominate the positive charge

making the ion-acoustic velocity very small

The FB instability can “continuously” heat the chromosphere

Magnetic field

Velocity

The electrons Hall drift produce the “electrostatic” Farley-Buneman instability that probably dissipates energy in the chromosphere

Convective motions should produce weak electric fields (~5 V/m) and drive the FB instability. Similar to the Earth ionosphere but in the Sun the instability is stronger and most everywhere because convective overshoot motions above granulation are above threshold most times.

Hall drift

Particle magnetization and FB instability threshold

New vs. old Model Set

0.1 1 10 100 1 103

1 104

1 105

1 106

4000

5000

6000

7000

8000

9000

CEFHPSR

Pressure (dyne cm^-2)

Tem

pera

ture

(K)

New semi-empirical chromospheric model set is being developed to match the CO lines and many others that the old models did not match. The old set of models needs update to match several lines, including CO.

0.01 0.1 1 10 100 1 103

1 104

1 105

1 106

4000

6000

8000

1001100210031004PSRR

Pressure (dyne cm^-2)T

empe

ratu

re (K

)

Revision to transition region (radiation/conduction+diffusion+flows)

2300 2350 2400 2450 2500 2550 26001 10

4

1 103

0.01

0.1

1

1 104

1 105

1 106

Rad. lossesTemperature

Height (km)

Rad

. los

ses

(erg

s^-

1 cm

^-3)

Tem

pera

ture

(K)

Energy balance transition regionstructure computed as in FAL. Optically thick and optically thin losses are included. Shown are the 306 model scaled with the usual (ne*nh)-1. Particle energy flux includes conduction and diffusion. TR is major energy sinkfor the corona and contributor to the UV radiation flux. Atomic data is being revised using CHIANTI

1 104

1 105

1 106

5 105

1 106

1.5 106

2500

3000

3500

4000

Energy FluxHeight

Temperature (K)

Ene

rgy

flux

(erg

s-̂1

cm

^-2)

Hei

ght (

km)

1 104

1 105

1 106

1 1025

1 1024

1 1023

1 1022

1 1021

2500

3000

3500

4000

Rad. lossesMech. heatingHeight

Temperature (K)

Scal

ed e

nerg

y lo

sses

/dis

sip

Hei

ght (

km)

Corona (radiation/conduction+wind+heating)

• Several magnetic field extrapolation methods produce more or less the field structure inferred from observed loops.

• Magnetic field extrapolations tend to fill the corona, but the emissions do not. Partial filling is necessary.

• Solar wind needs to be included for coronal holes.• Emission can be computed directly from loops and

wind models, but needs 3D and full Sun.• Coronal emission incident on the chromosphere has

some effects, especially on He spectrum.• For this task we intend to collaborate with groups

working on coronal loops and solar wind modeling.

Evaluating irradiance using disk masks

Using daily images of the solar disk various components are identified and a “mask” is produced. Daily spectra are computed using the semi-empirical models for the components (currently 7 components, will need 10). Comparison with SORCE data is shown for a few wavelengths (Lyα, 430 nm, and 656 nm).

SSI issues by SRPM

• Current research issues:– Discretization of continuous intensity distribution– UV & EUV surface features spectra distribution– Update plage & network chromospheric models– Inclusion of coronal holes and coronal loops– Status of magneto-convection simulations– 3D effects especially near the limb– Contributions to TSI variation by various bands– Spectral changes effects on Earth’s atmosphere

Courtesy of D. Braun

EARTH

Courtesy of D. Braun

EARTH

EARTH

10 20 30 40 50 60 70 80 900.006

0.0065

0.007

0.0075

Current rotationShifted previous rotation

Days since 2005/8/1

Ly

alp

ha i

rrad

ian

ce

Assuming previous curve is bad

10 20 30 40 50 60 70 80 900.006

0.0065

0.007

0.0075

Current rotationShifted previous rotation

Days since 2005/8/1

Ly

alp

ha i

rrad

ian

ce

Assuming previous curve is bad Images of the near-side produce daily masks of features

Using atmospheric models the spectrum is computed for any day

0.01 0.1 1 10 100 1 103

1 104

1 105

4

5

6 CEFH

Pressure (dyne cm^-2)

Log

(T)

1215.5 12161 10

3

1 104

1 105

1 106

1 107

CEFH

Wavelength

Inte

nsit

y (e

rg c

m^-

2 s^

-1 s

r^-1

)

0.01 0.1 1 10 100 1 103

1 104

1 105

4

5

6 CEFH

Pressure (dyne cm^-2)

Log

(T)

1215.5 12161 10

3

1 104

1 105

1 106

1 107

CEFH

Wavelength

Inte

nsit

y (e

rg c

m^-

2 s^

-1 s

r^-1

)

Without refinement the synoptic mask features obsolescence makes it bad

Synoptic masks are refined by applying trends and far-side imaging:

NOAA 10808 (far side)

NOAA 10808 (near side)

NOAA 10808 (far side)

NOAA 10808 (near side)

AR helioseismic image

AR backscattered image

Tools for forecasting solar irradiance